diff --git "a/9701.json" "b/9701.json" new file mode 100644--- /dev/null +++ "b/9701.json" @@ -0,0 +1,902 @@ +{ + "9701/astro-ph9701020_arXiv.txt": { + "abstract": "We examine the effects on the Fundamental Plane (FP) of structural departures from an $R^{1/4}$ galaxy light profile. We also explore the use of spatial (i.e.\\ volumetric) as well as projected galaxy parameters. We fit the Sersic $R^{1/n}$ law to the V-band light profiles of 26 E/S0 Virgo galaxies, where $n$ is a shape parameter that allows for structural differences amongst the profiles. The galaxy light profiles show a trend of systematic departures from a de Vaucouleurs $R^{1/4}$ law, in the sense that $n$ increases with increasing effective half-light radius $R_{\\rm e}$. This results in $R_{\\rm e}$, and the associated mean surface brightness within this radius, having systematic biases when constructed using an $R^{1/4}$ law. Adjustments to the measured velocity dispersion are also made, based upon the theoretical velocity dispersion profile shapes of the different $R^{1/n}$ light profiles, constructed assuming spherical symmetry and isotropic pressure support. We construct the FP for the case when structural homology is assumed (specifically, an $R^{1/4}$ law is fitted to all galaxies) and central velocity dispersions, $\\sigma _{0}$, are used. The plane we obtain is $R_{\\rm e}\\propto \\sigma _{0}^{1.33\\pm0.10}\\Sigma _{\\rm e}^{-0.79\\pm0.11}$, where $\\Sigma _{\\rm e}$ is the mean surface brightness within the projected effective radius $R_{\\rm e}$. This agrees with the FP obtained by others and departs from the virial theorem expectation $R\\propto \\sigma ^{2}\\Sigma ^{-1}$. We find that allowing for broken structural homology through fitting $R^{1/n}$ profiles (with $n$ a free parameter), but still using central velocity dispersions, actually increases the departure of the observed FP from the virial plane --- the increase in effective radius with galaxy luminosity (and $n$) is over-balanced by an associated decrease in the mean surface brightness. In examining the use of spatial quantities and allowing for the different velocity dispersion profiles corresponding to the observed light profiles, we find that use of the spatial velocity dispersion at the spatial half-light radius decreased the departure of the observed FP from the virial plane. (Use of the spatial half-light radius and mean surface brightness term had no effect on the FP as they are constant multiples of their projected values). Through use of the Jeans hydrodynamical equation, we convert the projected central aperture velocity dispersion, $\\sigma _{0}$, into the infinite aperture velocity dispersion, $\\sigma _{\\rm tot,n}$ (which is equal to one-third of the virial velocity dispersion). Using both the $R^{1/n}$ fit and $\\sigma _{\\rm tot,n}$ we obtain $R_{\\rm e,n}\\propto \\sigma _{\\rm tot,n}^{1.44\\pm0.11}\\Sigma _{\\rm e,n}^{-0.93\\pm0.08}$. Making the fullest allowance for broken structural homology thus brings the observed FP closer to the virial plane, with the exponent of the surface brightness term consistent with the virial expectation. ", + "introduction": "It has been known for some years that elliptical galaxies are well represented by a two dimensional surface in the space defined by their observable parameters \\cite{BaL83,Dj87a,Lyn8a}. This two dimensional manifold, commonly shown in the logarithmic 3-space of radius ($R_{\\rm e}$), intensity ($I_{\\rm e}$), and central velocity dispersion ($\\sigma _{0}$), has come to be known as the Fundamental Plane (FP) \\cite{DaD87,7Sb87}. The existence of such a plane has implications for the formation and evolutionary processes of elliptical galaxies. Furthermore, the small scatter about the FP makes it a useful tool for estimating distances to elliptical galaxies. Acknowledging some degree of scatter about the FP, either intrinsic to the galaxies and/or due to measurement errors, the deviations from the FP have been interpreted as peculiar velocities \\cite{7Sa87,DaF90}. For such methods of studying peculiar velocity fields and bulk flow motions, the exact FP used is a crucial factor. Under the assumptions of kinematic and structural homology (i.e. identical form and shape for the kinematic and density profiles) amongst galaxies, the virial theorem predicts that galaxies will occupy a plane in log-space amongst the observables radius ($R_{\\rm e}$), mass surface density ($\\eta _{\\rm e}$), and central velocity dispersion ($\\sigma _{0}$) \\cite{F7S87}. The transformation between this plane and the observed plane is given by the mass-to-light ratios, M/L, of the galaxies. Furthermore, the tight constraints on the thickness of the observed plane and its negligible curvature, implies that M/L is a power law function of the observables. The mass-to-light ratio has important implications for galaxy formation \\cite{F7S87,Dj87b} and so one would like to be able to test the assumptions of homology that go into the model. The observed difference between the FP and the relation predicted from the virial theorem is usually viewed as a systematic variation in the mass-to-light ratio along the FP, given by $M/L\\sim L^{\\alpha }$, where $\\alpha \\sim 0.25\\pm 0.05$ \\cite{FaJ76,F7S87,Dj87b}, and it has been shown to vary with bandpass (Djorgovski \\& Santiago 1993; Pahre, Djorgovski \\& de Carvalho 1995). Many causes of such an effect have been explored and dismissed to varying extents. Possible stellar population differences along the FP \\cite{Dj87b,DaS93,RaC93} seem unlikely due to the required fine-tuning of the IMF or $M/L$ required in order to maintain the thinness of the observed FP. A systematic variation in stellar age along the plane could produce the observed tilt, but would require a large conspiracy in the formation of all galaxies of a given mass in order to be consistent with the observed small thickness of the FP. To examine stellar population effects, Pahre et al.\\ \\shortcite{PDC95} tested the significance of the mass-metallicity relation \\cite{Guz02} along the FP, by constructing a near-infrared FP. Sampling the older distribution of stars, their K-band photometry is less sensitive than optical wave-bands to differences in galaxy metallicity which appear as line-blanketing effects in the optical \\cite{F7S87}. The near-infrared FP obtained did not differ markedly from the optical FP, with the deviation found being attributed to the reduction in metallicity effects between galaxies. This agrees with earlier work by Dressler et al.\\ (1987b), Recillas-Cruz et al.\\ \\shortcite{Rec90}, and Djorgovski \\& Santiago \\shortcite{DaS93}, where the metallicity effect was shown to be insufficient to explain the entire observed FP departure from the virial plane. This is also shown to be the case with current population synthesis models, where metallicity effects result in only a small tilt in the plane \\cite{Wor94}. This has led to investigations that test the assumptions of kinematic and structural homology and measure the influence of such a breakdown upon the FP (Ciotti, Lanzoni \\& Renzini 1996; Hjorth \\& Madsen 1995). Pahre et al.\\ \\shortcite{PDC95} also concluded that systematic departures from structural and dynamical homology amongst the galaxies may be responsible for the departure between the observed and the virial form of the FP. Capelato, de Carvalho \\& Carlberg \\shortcite{CCC95}, using computer model simulations of merged galaxies, claim that the slope of the FP is largely explained by broken homology in the velocity distribution. However, Ciotti et al.\\ \\shortcite{CLR96} and Ciotti \\& Lanzoni \\shortcite{CaL96}, using spherical models with the velocity anisotropy described by the Osipkov-Merritt formula, find it is unlikely that orbital anisotropy plays a major role in producing the FP tilt. Anisotropy in the velocity distribution has also been rejected by Djorgovski \\& Santiago (1993) as the sole cause for the $M/L\\sim L^{\\alpha }$ relation. Ciotti et al.\\ \\shortcite{CLR96} suggested that a systematic trend in the shape of the elliptical galaxy light profiles seems to be the best explanation. This possibility has been entertained before (Djorgovksi, de Carvalho, \\& Han 1988; Djorgovski \\& Santiago 1993), but differing profile shapes causing the tilt of the FP has not been explored with observational data. Hjorth \\& Madsen (1995) parameterised broken structural homology in their models of elliptical galaxies based on the statistical mechanics of violent relaxation \\cite{HaM91}. They found that for a given correlation between the galaxy structure and its luminosity, such non-homology is not inconsistent with the observed FP having a constant $M/L$. Unfortunately, their data set did not enable them to draw firm conclusions. There is, however, a growing wave of data supporting the existence of broken structural homology, amongst dwarf galaxies \\cite{Dav88,BaC91}, spiral bulges \\cite{Cap87,Cap89,APB95,CJB96}, elliptical galaxies (Michard 1985; Caon, Capaccioli \\& D'Onofrio 1993; Burkert 1993; Hjorth \\& Madsen 1995), and brightest cluster galaxies \\cite{Sch86,Gra96}. In this paper, using a set of observed surface brightness profiles and measured velocity dispersions, we explore the influence of broken structural and dynamical homology upon the FP. Using the Sersic $R^{1/n}$ light profile \\cite{Ser68}, the generalised form of the de Vaucouleurs $R^{1/4}$ law (de Vaucouleurs 1948, 1953) with $n$ a free parameter, we can allow for structural differences between galaxies. The $R^{1/4}$ law has two parameters which act as physical scales, scaling the radius and the surface brightness and leaving the profile shape fixed. The $R^{1/n}$ profile has the additional parameter, n, which permits different shapes for the galaxy light profile and therefore allows us to examine broken structural homology amongst galaxies. The Sersic law can take the form of both an exponential disk ($n$=1) and the de Vaucouleurs light profile ($n$=4). It can also approximate a power-law and provides intermediate forms between these common profile shapes by varying its shape parameter $n$. Caon et al.\\ \\shortcite{CCDet} and Graham et al.\\ \\shortcite{Gra96}, and references therein, have shown that a range in galaxy profile shapes does exist, and more importantly that a systematic trend between profile shape and galaxy size exists such that the larger galaxies have less curvature in their light profiles than the smaller galaxies. In this paper we examine the effect of this on the FP. We also explore the use of volumetric galaxy parameters as well as using the standard projected quantities. It is possible to obtain the three dimensional structural quantities by deprojecting the $R^{1/n}$ light profiles (Ciotti 1991). The volumetric kinematical quantity is derived from application of the Jeans hydrodynamical equation to the $R^{1/n}$ model, calibrated by the projected central velocity dispersion. In section 2 we describe our sample of elliptical galaxies and the derivation of their structural parameters. An overview of the parameters used in the construction of the FP and their relation to the ideal virial plane parameters is given in section 3. We use multivariate statistics to describe and construct the FP in section 4. Our results are discussed in section 5, and our conclusions presented in section 6. ", + "conclusions": "A range in profile shapes is found to exist for the sample of 26 elliptical galaxies used in this study. The shape parameter $n$ is found to have a physical association with galaxy size and mass, in the sense that a larger value of $n$ corresponds to a larger galaxy. Allowances for different profile shapes and the use of infinite aperture velocity dispersion measurements, which allow for different velocity dispersion profile shapes, results in a modification to the FP. While more clusters need be studied, our work shows that broken structural homology is not solely responsible for the tilt of the FP, but allowing for variations caused by such structural differences does reduce the tilt. It is likely that the departure of the observed FP from the virial plane will be further reduced by using infrared luminosity profile measurements to minimise stellar population effects. Our best estimate of the FP using the parameters from an $R^{1/n}$ profile and the total aperture velocity dispersion is $R\\propto \\sigma ^{1.44\\pm0.11}\\Sigma ^{-0.93\\pm0.08}$. The exponent of the surface brightness term is consistent with the prediction of the virial theorem. \\subsection*" + }, + "9701/astro-ph9701111.txt": { + "abstract": "We use Bayesian methods to analyze the distribution of gamma ray burst intensities reported in the {\\it Third BATSE Catalog} (3B catalog) of gamma ray bursts, presuming the distribution of burst sources (``bursters'') is isotropic. We study both phenomenological and cosmological source distribution models, using Bayes's theorem both to infer unknown parameters in the models, and to compare rival models. We analyze the distribution of the time-averaged peak photon number flux, $\\F$, measured on both 64~ms and 1024~ms time scales, performing the analysis of data based on each time scale independently. Several of our findings differ from those of previous analyses that modeled burst detection less completely. In particular, we find that the width of the intrinsic luminosity function for bursters is unconstrained, and the luminosity function of the actually observed bursts can be extremely broad, in contrast to the findings of all previous studies. Useful constraints probably require observation of bursts significantly fainter than those visible to BATSE. We also find that the 3B peak flux data do not usefully constrain the redshifts of burst sources; useful constraints require the analysis of data beyond that in the 3B catalog (such as burst time histories), or data from brighter bursts than have been seen by BATSE (such as those observed by the {\\it Pioneer Venus Orbiter}). In addition, we find that an accurate understanding of the peak flux distributions reported in the 3B almost certainly requires consideration of data on the temporal and spectral properties of bursts beyond that reported in the 3B catalog, and more sophisticated modeling than has so far been attempted. We first analyze purely phenomenological power law and broken power law models for the distribution of observed peak fluxes. We find that the 64~ms data is adequately fit by a single power law, but that the 1024~ms data significantly favor models with a sharp, steep break near the highest observed fluxes. At fluxes below the break, the distribution of 1024~ms fluxes is flatter than that of 64~ms fluxes. Neither data set is consistent with the power law distribution expected from a homogeneous, Euclidean distribution of sources. Next we analyze three simple cosmological models for burst sources: standard candles with constant burst rate per comoving volume; a distribution of standard candle sources with comoving burst rate proportional to a power law in $(1+z)$, and a bounded power-law burster luminosity function with constant comoving burst rate but variable power-law index and luminosity bounds. We find that the 3B data can usefully constrain the luminosity of a standard candle cosmological population of bursts if there is no density evolution. But the 3B data allow strong density evolution and arbitrarily broad luminosity functions; consequently, they do not usefully constrain the redshifts or luminosities of cosmological burst sources. We elucidate the properties of the models responsible for these results. For sufficiently flexible models, the inferred values for parameters describing the shapes of the distributions of 64~ms and 1024~ms peak fluxes formally differ at the 68\\%--95\\% level. Because the measurements on these two timescales are not independent, it is difficult to ascertain the true significance of this discrepancy; since many bursts are common to both data sets, it is likely its significance is larger than these formal values indicate. In addition, the inferred amplitude (in bursts per year) of the distribution of 64~ms peak fluxes is about twice that of 1024~ms peak fluxes. These results strongly suggest that a complete understanding of the measured peak flux distributions requires simultaneous modeling and analysis of temporal properties of bursts. We study models that attempt to reconcile the two data sets by accounting for ``peak dilution,'' the underestimation of the peak intensity that results from using data accumulated over a timescale exceeding the peak duration. A phenomenological model strongly correlating peak duration with peak flux is moderately successful at reconciling the data. A model that correlates peak duration with peak flux due to cosmological time dilation and relativistic beaming is less successful, but remains of interest in that it is a simple physical model illustrating how one can jointly model and analyze temporal and spectral properties of bursts with peak flux data. A more rigorous accounting for the differences between the 64~ms and 1024~ms data requires analysis of temporal and spectral information about bursts beyond that available in the 3B catalog. ", + "introduction": "In the absence of direct measurement of the distances to burst sources (``bursters''), or association of bursts with well-localized counterparts, we must infer the spatial and energy distribution of bursters from the observed distribution of burst strengths and directions. The complexity of the burst data makes this task considerably more difficult than it might at first appear because we must account for several subtle biases and selection effects. So far, analyses of burst data have relied on largely ad hoc choices of statistics designed to circumvent some of these effects. But differing choices of statistic or analysis method by different investigators have led to some controversy over the implications of the burst data. Of particular relevance to this work are the varying conclusions of numerous previous studies of isotropic models for the burst source distribution. On a purely phenomenological level, investigators differ over whether the logarithmic slope of the distribution of burst intensities exhibits a significant change in slope (cf.\\ Loredo and Wasserman 1993; Wijers and Lubin 1993; Petrosian, Azzam, and Efron 1994). In the context of cosmological models, investigators have reached a variety of inconsistent conclusions about possible characteristics of the burst source distribution, including the value and uncertainty for the luminosity of standard candle sources (cf.\\ Dermer 1992; Loredo and Wasserman 1993), and the ability of the data to constrain the width of the luminosity function of bursters (cf.\\ Loredo and Wasserman 1993; Horack, Emslie, and Meegan 1994; and Cohen and Piran 1995) or the redshift of the faintest bursters (cf.\\ Loredo and Wasserman 1993; Emslie and Horack 1994; Cohen and Piran 1995). These differences result largely from methodological differences among the published studies. Without a vastly larger dataset, only careful attention to methodological issues can identify the correct inferences. In the first paper of this series (Loredo and Wasserman 1995; hereafter LW95), we described the Bayesian methodology for inferring the spatial and energy distribution of burst sources. Instead of constructing a customized statistic in an attempt to circumvent the various biases and selection effects that might enter inferences, we start from simple models (based on the Poisson distribution) for burst occurence and detection, and directly calculate the probability for the observed data: the likelihood function. The likelihood function describes how well a model can account for the {\\it joint differential distribution} of observed burst strengths and directions, and accounts for biases and selection effects by construction, rather than trying to circumvent them by a clever choice of statistic. Indeed, from the Bayesian point of view, there is no freedom of choice regarding what statistic to use and how to use it; the data enter Bayes's theorem in the likelihood function, and the rules of probability theory dictate both how to calculate the likelihood function and how to manipulate it to make inferences. This methodology offers several advantages over rival methods: (1) it does not destroy information by binning or averaging the data (as do, say, $\\chi^2$, $\\langle V/V_{\\rm max}\\rangle$, and analyses of flux or angular moments); (2) it straightforwardly handles uncertainties in the measured quantities; (3) it analyzes the strength and direction information jointly; (4) it uses information available about nondetections; and (5) it automatically identifies and accounts for biases and selection effects, given a precise description of the experiment. In this work we use the Bayesian methodology to make inferences about isotropic models for the distribution of burst sites, using data from the {\\it Third BATSE Catalog} (Fishman et al.\\ 1996, hereafter the 3B catalog; the 3B catalog inherits some properties of the First BATSE (1B) catalog described in Fishman et al.\\ 1994). A companion paper (Loredo and Wasserman 1996) uses the same methodology to make inferences about anisotropic models, including comparisons of isotropic and anisotropic models. In the next section, we briefly review the Bayesian methodology and the form of the likelihood function for burst data. The 3B catalog does not provide all of the information required for a rigorous analysis, so we are forced to approximate some of the quantities required for calculating the likelihood. In \\S~3 we describe the approximations necessary to calculate the likelihood function based on the data available in the 3B catalog. Consistency requires that data for some bursts be omitted from the analysis when the approximations fail for those bursts, so we carefully discuss selection of the analyzable data. The most serious data cut arises from inaccuracy of the approximation used to calculate the detection efficiency reported in the 1B catalog for bursts with low fluxes. We describe extensive simulations we performed to quantify the inaccuracy of the approximation, and find that a significant fraction of bursts must be omitted from the analysis to avoid seriously corrupting the inferences drawn. Previous analyses of the BATSE data have not worried about inaccuracies introduced by improperly including dim bursts; this may account for some of the differences between our conclusions and those reached in other analyses. We carry out our analysis in \\S~3 and throughout the remainder of the paper using data based on peak photon number flux measurements taken on both 64~ms and 1024~ms time scales. We analyze these data separately; they are not independent, but their dependence is too complicated to quantify with the data tabulated in the 3B catalog alone. An important conclusion of our analyses is that inferences drawn from these two data sets differ substantially, although the formal significance of the discrepancy cannot be determined without more information about the bursts comprising the catalog. We make an effort to understand and quantify this discrepancy in the final sections of this paper. Although we cannot conclusively identify the reason for the discrepancy, we suggest that it is a result of burst light curves having peak durations that are often significantly shorter than 1024~ms (and possibly shorter than 64~ms) and that may be correlated with burst intensity. The resulting errors in peak flux estimates distort the peak flux distribution; the distortion differs for the two data sets, and can potentially account for their significantly different shapes. We begin, however, with models that do not attempt to account for any effects the different time scales of the data sets may have on the shape of the observed distribution of burst intensities. For brevity, we deem such models ``simple'' models. In \\S~4 we present results of analyses of simple phenomenological models for the differential burst rate (the burst rate per unit peak flux, $\\F$). These models help us ascertain what features must be present in the burst rate without committing us to a particular physical explanation for these features. They indicate that the distribution of 64~ms peak fluxes is adequately fit by a single power law whose logarithmic slope is very significantly different from the $-2.5$ value associated with the differential rate for an unbounded homogeneous population of sources. There is no significant evidence for steepening of the distribution with intensity, contrary to thesuggestions of such steepening we found in the 1B catalog (Loredo and Wasserman 1993). In contrast, the 1024~ms peak flux data prefer models with a broken power law distribution, and the low flux part of the distribution is significantly shallower than the distribution of 64~ms peak fluxes. In \\S\\ 5 we present results of analyses of three simple cosmological models for the burst source distribution. The simplest model presumes that all burst sources have the same intrinsic luminosity (they are ``standard candles''), and that the burst rate per unit comoving volume is constant with redshift, $z$. Next, we consider standard candle sources with a burst rate density that varies as a power of $(1+z)$. Finally, we consider models with power-law luminosity functions and constant comoving burst rate density. We find that our ignorance of the additional parameters in models with density evolution or a luminosity function greatly weakens our ability to learn about the spatial distribution and luminosity of the sources of bursts. As with the simple phenomenological models, these models reveal systematic discrepancies between the 64~ms and 1024~ms data. In \\S~6 we discuss how properties of burst light curves might lead to time scale-dependent distortions of the observed flux distribution qualitatively capable of reconciling the two data sets. That such effects might prove crucial for making inferences from burst peak flux data was already anticipated in LW95; other authors have also previously remarked on the importance of these effects for understanding the flux distribution of bursts (Lamb, Graziani, and Smith 1993; Petrosian, Lee, and Azzam 1994). Following a general discussion, we analyze a more complicated phenomenological model than those of earlier sections that takes into account the different measuring time scales of the data sets. It is only moderately successful at reconciling the data sets, but remains of interest as an example of how one can explicitly account for time scale dependent effects in models. In \\S~7 we analyze a final physical model that draws together several of the lines of thought developed in the preceding sections. In this model, burst sources are standard candles and standard clocks in their rest frames, but undergo relativistic motion with respect to locally comoving observers. An isotropic distribution of beaming angles with respect to the line of sight results in an effective luminosity function that is a power law if the rest frame emission spectrum is a power law. In the absence of time scale effects, this model is thus identical to the cosmological model with a power law luminosity function considered in \\S~5, except that the power law index is a function of the burst spectral index, rather than a free parameter. Beaming also results in bursts having observed peak durations that are a function of luminosity, allowing us to model time scale effects; cosmological redshift additionally correlates duration and peak flux. Thus besides being a model of intrinsic physical interest, this model indicates how both the temporal and spectral properties of bursts can influence the flux distribution in a manner that can be straightforwardly modeled. The final section summarizes our findings and their implications. Two technical appendices describe the details of our cosmological models. %=============================================================================== ", + "conclusions": "We have analyzed the 64~ms and 1024~ms peak flux data in the 3B catalog using the Bayesian method described in detail by Loredo and Wasserman (1995). The method identifies several shortcomings of the summaries of the data comprising the 3B catalog that prevent consistent analyses of the entire catalog. In particular, counting uncertainties and atmospheric scattering were omitted from the calculation of the detection efficiencies reported in the 3B catalog, requiring that the dimmest 38\\% of the 64~ms bursts, and the dimmest 16\\% of 1024~ms bursts, be omitted from any analysis of the peak flux distribution. We have used the resulting self-consistent data sets to analyze a variety of phenomenological and physical (cosmological) models for burst sources that presume burst sites are distributed isotropically. A companion paper presents analyses of anisotropic models that associate some or all bursts with an extended Galactic halo. %------------------------------------------------------------------------------- \\subsection{Simple Phenomenological Models} Our analysis of phenomenological models based on power laws and broken power laws verifies that neither the 64~ms nor the 1024~ms data is consistent with a homogeneous (Euclidean) distribution of sources, for which the differential burst rate obeys $dR/d\\F \\propto \\F^{-5/2}$. There is no significant evidence for a break in the logarithmic slope of the distribution of 64~ms peak fluxes, but there is moderately significant evidence for such a break to the homogeneous $\\gamma=2.5$ slope in the 1024~ms data, and stronger evidence for a steep cutoff in the distribution of bursts with $\\F_{1024} \\gtrsim 40$~cm$^{-2}$~s$^{-1}$. Also, the power law indices that best describe the low flux portion of each data set differ with at least moderate significance, the low-flux distribution of 1024~ms peak fluxes being somewhat flatter than those favored for the 64~ms peak fluxes. The different inferred shapes of the two data sets are not simply due to their different sizes and dynamic ranges. This argues that a full understanding of the shape of the observed peak flux distribution requires explicit consideration of the temporal structure of burst light curves. A simplified analysis (summarized below) indicates that the structure in the 1024~ms flux distribution is an artifact of its longer measurement time scale, so that the shape of the 64~ms flux distribution is more representative of the shape of the distribution of instantaneous peak fluxes. Several physical models for the distribution of burst sources in space and luminosity predict flux distributions that are well approximated by power laws and broken power laws, as noted in \\S~4. Quite generically, information about characteristic length and luminosity scales in such models is revealed by a change in the logarithmic slope of the flux distribution from a relatively flat differential distribution for dim bursts to a steeper $\\F^{-5/2}$ distribution for bright bursts. That there is no evidence for such a change in the 64~ms data, and only moderate evidence for a change to a $\\F^{-5/2}$ power law in the 1024~ms data, presaged the conclusions we found in our analyses of physical models: the data are unable to constrain properties of cosmological populations of burst sources. %------------------------------------------------------------------------------- \\subsection{Simple Cosmological Models} We analyzed three simple cosmological models in an effort to determine whether the 3B data could detect or rule out evolution of the burst rate density with redshift, and whether the data could constrain the width of the burster luminosity function. The data are unable to discriminate among homogeneous standard candle models and models with strong density evolution or broad luminosity functions. As a consequence, the luminosity of burst sources is uncertain over many orders of magnitude, and the typical redshifts of observed bursts can be as small as a few tenths or $\\gtrsim 20$. The upper limit could almost certainly be reduced by considering PVO data, since it depends on locating the flux where the flux distribution steepens to $dR/d\\F \\propto \\F^{-5/2}$. A stronger constraint may arise from the absence of large time dilation in burst lightcurves, although the wide variety of temporal behavior exhibited by bursts severely complicates the modeling and detection of such time dilation (see, e.g., Mitrofanov 1996, Fenimore 1996, and Norris and Nemiroff 1996, who reach conflicting conclusions regarding the presence of ``time stretching'' in BATSE data). A lower limit on the redshift of observed sources may be sought more effectively from the absence of anisotropy in the distribution of burst directions (as would appear if many bursts were visible from within the local supercluster, for example) than from the flux distribution (see, e.g., Quashnock 1996). Since such anisotropy should correlate with burst intensity, our Bayesian methodology is an ideal tool for rigorously studying it. Even in the absence of strong density evolution (in which case the observed bursts have typical redshifts of a few tenths), the width of the luminosity function for burst sources is unconstrained and could span several orders of magnitude. Unfortunately, we find that the uncertainty in the luminosity of the brightest burst sources is comparable in size to the range of the luminosity function, and thus is largely unconstrained by the BATSE data. The addition of PVO data is unlikely to strengthen this constraint, because models with luminosity functions of very different widths distinguish themselves at low fluxes rather than at large fluxes. Without a vastly larger data set, the best hope for constraining the width of the luminosity function of cosmological burst sources is to obtain data on the distribution of bursts with fluxes well below the threshold of the BATSE detectors. \\clearpage %------------------------------------------------------------------------------- \\subsection{Consideration of Temporal\\\\ and Spectral Properties of Bursts} Analyses with sufficiently flexible models reveal systematic differences between the shapes of the distributions of 64~ms and 1024~ms peak fluxes. Parameters inferred from the two data sets differ with moderate significance presuming they are independent. But since the two data sets are not independent (well over half of the 1024~ms bursts triggered on the 64~ms timescale) one would expect close agreement between the inferred values; the discrepancy between the inferences is thus probably very significant. In addition, the normalizations of the two distributions are extremely different, the number of 1024~ms bursts with fluxes above 1.5~cm$^{-2}$~s$^{-1}$ (the cutoff for 64~ms bursts) being only 56\\% of the number of 64~ms bursts. In LW95 we argued that explicit consideration of the temporal properties of bursts would be necessary for understanding the distribution of measured burst fluxes. The disparity between the 64~ms and 1024~ms data supports this claim. We therefore attempted to model the data in a manner that crudely accounts for ``peak dilution'': the underestimation of peak flux that occurs when estimating peak flux with data from a time interval longer than the peak duration. If peak duration is correlated with peak flux, peak dilution can result in an observed peak flux distribution that is different in shape from the underlying actual peak flux distribution. Unfortunately, the 3B catalog contains no direct information about the peak durations of bursts. Thus we have been able to perform only illustrative calculations that show how explicit consideration of temporal properties of bursts might enter an analysis of the flux distribution. Should peak duration measurements become available, more reliable and definitive analyses will be possible. We analyzed a purely phenomenological model in which bright bursts were presumed to have shorter peak durations than dim bursts (the qualitative behavior expected in cosmological models). Thus bright bursts have their peak fluxes systematically underestimated when long measuring time scales are used, steepening the flux distribution at bright fluxes. This model is moderately successful in reconciling the shapes of the two data sets, and in particular is capable of producing a cutoff in the distribution of observed peak fluxes, as is seen in the 1024~ms data. Additionally, we analyzed a physical model in which a cosmological population of relativistically beamed sources that are standard candles and clocks produces an apparent distribution of sources with a broad luminosity function and distribution of peak durations, due to the distribution of the angle between the source velocity and the line of sight. Besides correlating burst duration and peak flux, this model also correlates the burst spectrum with peak flux and duration. However, it does not successfully account for the differences between the two data sets. Nevertheless, it is of intrinsic physical interest, and further, it is the simplest model illustrating how the spectral and temporal properties of bursts can enter the analysis of the flux distribution. We thus conclude that the BATSE peak flux data cannot usefully constrain cosmological models for burst sources. Useful constraints from peak fluxes alone will result only from consideration of data about the infrequent bright bursts that BATSE has not yet seen (to constrain the redshifts of the most distant observed sources), and about bursts dimmer than BATSE is capable of seeing (to constrain the width of the luminosity function). Joint analyses of temporal and spectral properties of bursts with their peak fluxes may well provide more useful constraints on cosmological models. The Bayesian methodology adopted here is the ideal tool for such a joint analysis. %===============================================================================" + }, + "9701/astro-ph9701093_arXiv.txt": { + "abstract": "Recent evidence on the metal content of the high-redshift \\lya forest seen in quasar spectra suggests that an early generation of galaxies enriched the intergalactic medium (IGM) at $z\\gtrsim 5$. We calculate the number of supernovae that need to have taken place to produce the observed metallicity. The progenitor stars of the supernovae should have emitted $\\sim 20$ ionizing photons for each baryon in the universe, i.e., more than enough to ionize the IGM. We calculate that the rate of these supernovae is such that about one of them should be observable at any time per square arc minute. Their fluxes are, of course, extremely faint: at $z=5$, the peak magnitude should be $K=27$ with a duration of $\\sim$ 1 year. However, these supernovae should still be the brightest objects in the universe beyond some redshift, because the earliest galaxies should form before quasars and they should have very low mass, so their luminosities should be much lower than that of a supernova. We also show that, under the assumption of a standard initial mass function, a significant fraction of the stars in the Galactic halo should have formed in the early galaxies that reionized and enriched the IGM, and which later must have merged with our Galaxy. These stars should have a more extended radial distribution than the observed halo stars. ", + "introduction": "Ever since the discovery of the first high-redshift quasar (Schmidt 1965), quasars have maintained their title as the objects with the highest known redshift; the present record holder is a quasar at $z=4.89$ (Schmidt, Schneider, \\& Gunn 1991). Nevertheless, the highest known redshifts of galaxies have followed closely behind, with bright radio galaxies having been found up to $z=4.45$ (Lacy \\etal 1994; Rawlings et al.\\ 1996); more recently, galaxies with high star formation rates have been identified from interstellar absorption lines at $z=2$ to 3 (Steidel et al.\\ 1996) and from the \\lya emission line at $z=4.55$ (Hu \\& McMahon 1996). In fact, if quasars are related to supermassive black holes that formed in the centers of high-redshift galaxies, we should expect that many galaxies already existed before the first quasars appeared. In any `bottom-up' theory where the observed structure in the universe forms by hierarchical gravitational collapse, and the primordial density fluctuations extend to sufficiently small scales, the first galaxies to form must have had much smaller masses than the present galaxies. The first stars should have formed in systems with velocity dispersions of $\\sim 10 \\kms$ or lower, corresponding to the lowest temperatures ($T\\sim 10^4$ K) that allow cooling and dissipation of the gas by atomic processes (systems with even lower virial temperatures can cool and dissipate through molecular hydrogen, but this cooling process should be suppressed by photodissociation of the molecules after emission of a number of UV photons that is much smaller than that needed to reionize the universe; see Haiman, Rees, \\& Loeb 1996). These systems would be very unlikely to form quasars, because even a small fraction of their baryons turning into stars should provide sufficient energy (via ionization, stellar winds or supernovae) to expel the remaining gas from the shallow potential well (e.g., Couchman \\& Rees 1986; Dekel \\& Silk 1986). Deeper potential wells, forming at later epochs, are probably needed to form supermassive black holes in galactic centers. Even if all the baryons were converted into stars very efficiently in these early dwarf galaxies, with a baryonic mass $M_b \\lesssim 10^8 \\msun$, their total stellar luminosity would be much smaller than in $L_*$ galaxies at present, simply due to their small mass. Since only a small fraction of the baryons in these systems is likely to turn into stars before the gas is ejected, the total stellar mass in the first galaxies to form in the universe should be much smaller than $10^8 \\msun$. This implies that a supernova in one of these first galaxies to form will be far brighter than the galaxy itself. Thus, the brightest probes of the era when the reionization of the intergalactic medium (IGM) started should be supernovae in very small galaxies, caused by the death of the same stars responsible for the first ionizing photons. In this paper, we shall estimate the number of supernovae that should have taken place in these galaxies and should be observable at very high redshift, and their apparent magnitudes. ", + "conclusions": "As the observational techniques improve our ability to detect extremely faint sources, and higher redshift objects can be searched for to continue unravelling the history of galaxy formation, supernovae should become the brightest observable sources. These supernovae created the heavy elements that were expelled to the IGM, and their progenitor stars are the most likely sources of the photons that reionized the universe. The expected rates of these supernovae, calculating under the assumption of a high baryon density ($\\Omega_b h^2 = 0.025$), and an average metal production of $\\bar Z = 10^{-2} \\zsun$, is as high as 1 supernova per square arc minute per year. To detect the supernovae, the flux limits of the faintest sources detectable with our telescopes will probably need to be pushed by another $\\sim 2$ magnitudes, although the first examples might be discovered at brighter fluxes behind clusters of galaxies, using the lensing magnification. Any low-mass stars that were formed in the first small galaxies where these supernovae took place should be observable today. We have argued that, if the IMF in these galaxies was similar to the present one in our galactic disk, the Population III stars are likely to account for a large fraction of the stars in our galactic halo, although most of them should be in an as yet undetected outer halo with a shallower density profile than the known, inner stellar halo." + }, + "9701/astro-ph9701087_arXiv.txt": { + "abstract": "We present phase-resolved spectroscopy of the eclipsing AM Herculis star UZ For obtained when the system was in its low state of accretion. Faint residual H$\\alpha$-emission and NaI absorption were used to trace the secondary star and infer its orbital velocity $K_2$. The measured radial velocity amplitude of NaI $K_2 * \\sin{i} = 285 \\pm 50$\\,\\kmps\\ suggests a low-mass white dwarf with $M_{\\rm wd} = 0.44 \\pm 0.15$\\,\\msun (1$\\sigma$-errors). The H$\\alpha$ emission line on the other hand, visible only for part of the orbital cycle and supposed to originate only on the illuminated hemisphere facing the white dwarf, displays a similar radial velocity amplitude, $K'_2 \\sin i = 308\\pm 27$\\,\\kmps. The standard $K_2$-correction applied by us then suggests a white dwarf mass of up to 1\\,\\msun. Compared with earlier results the new ones enlarge the window in which the white dwarf mass may lie and resolves the conflict between mass estimates based on photometry and spectroscopy. They leave some ambiguity in the location of emission and absorption components in these and former observations. ", + "introduction": "The eclipsing polar (AM Herculis binary) UZ For became of considerable interest soon after its discovery and optical identification (Osborne et al.~1988, Beuermann et al.~1988, Paper I) because of its high magnetic field, its period of 126.5\\,min, close to the lower edge of the CV period gap, and the likely high mass of the accreting white dwarf, $M_{\\rm wd} > 0.93$\\,\\msun\\/ (90\\% confidence). The high mass of the white dwarf was based on the observed high radial velocity amplitudes of H$\\alpha$ emission and NaI absorption lines originating at the secondary star. The white dwarf in \\uz\\ thus seemed to be more massive than those in other AM Her binaries and than single ones. On the assumption that \\uz\\ resumes accretion at the observed period after crossing the period gap (a quantity which is dependent on \\mwd), Hameury et al.~(1988) argued also for a massive white dwarf from evolutionary considerations, \\mwd\\ $\\simeq 1.2$\\,\\msun. High-speed photometry presented by Bailey and Cropper (1991) with resolved ingress and egress of the white dwarf, however, suggests that the white dwarf is a normal one with standard mass \\mwd = 0.61 -- 0.79 \\msun. ", + "conclusions": "Both radial velocity measurements, H$\\alpha$ and NaI, can be used to re-estimate the mass of the white dwarf. This is done using the following assumptions: (a) The secondary fills its Roche lobe completely and is (b) a main sequence star; (c) The mass-radius relation of the secondary is of the form $R_2 = \\alpha M_2^\\gamma$, and $R_2$ is set equal to the spherical equivalent Roche radius $R_s$; (d) $R_s$ is approximated by the formula given by Eggleton (1983) as a function of the mass ratio $Q$; (e) The measured radial velocity amplitude $K_{\\rm Na}$ represents the true orbital velocity of the companion star (apart from a $\\sin{i}$ factor); (f) The full eclipse length of the white dwarf is 463.9\\,sec corresponding to a half phase angle of 11.00\\degr. Allen et al.~(1989) derived a full eclipse length of the hot accretion spot on the white dwarf surface of $466.5 \\pm 2.5$\\,sec. This number has to be corrected for the offset between the white dwarf center and the spot. We assume for this correction a binary separation of $a = 5.30 \\times 10^{10}$\\,cm, a white-dwarf radius $R = 8.7 \\times 10^8$\\,cm, an azimuth of the spot of -45\\degr\\ and a colatitude of 150\\degr. \\begin{figure} \\psfig{figure=qi_rel_ps,width=87mm,bbllx=49mm,bblly=72mm,bburx=166mm,bbury=240mm,clip=} \\caption[meanlc]{\\label{qirel} Mass estimate of the white dwarf in \\uz. The three panels show from top to bottom the inclination $i$, the mass of the white dwarf and the predicted orbital velocity of the secondary star as a function of the mass ratio $Q$ for given period, eclipse length and adopted mass-radius relation for the secondary star, respectively. In each of the panels the three more or less parallel lines indicate predictions made by different mass-radius relations for the secondary star (solid: Neece 1984, dotted: Caillault \\& Patterson (1990), short dashed: VandenBerg et al.~1983). The solid line connecting the lower left with the upper right in the upper panel relates the eclipse length with $i$ and $Q$ and was calculated for an adopted eclipse length of 463.9 sec. The dots below this curve are drawn from Bailey \\& Cropper (1991) who used an unknown eclipse length. The three mainly vertical lines to the left in the upper panel are mass functions calculated for a velocity of 285\\,\\kmps\\/and different M/R-relation. The long-dashed line and the dashed-dotted line are mass functions for 385\\,\\kmps\\/and 416\\,\\kmps, respectively, using the M/R-relation given by Neece. } \\end{figure} The relation between inclination $i$ and mass ratio $Q$ for given eclipse length due to Chanan et al.~(1976) is shown as a solid curve in the upper panel of Fig.~\\ref{qirel} connecting the lower left with the upper right edge. This relation is purely geometrical and not dependent on e.g.~ZAMS mass-radius assumptions. It seems, however, that we used a different eclipse length of the white dwarf than Bailey \\& Cropper (1991) did in their analysis (small dots below the solid curve are depicted from their Table 2). Unfortunately they do not quote their measured eclipse length, so that a direct comparison is not possible. The highest inclination and highest mass ratio compatible with the observed eclipse length are $i_{\\rm max} = 86.2\\degr$ and $Q_{\\rm max} = 9.20$. The three parallel lines in the upper panel to the left are the mass functions calculated for the nominal velocity of 285\\,\\kmps using three different mass radius relations. The middle line is based on Neece (1984, N84), the dotted line on Caillault \\& Patterson (1990, CP90) and the short dashed line on VandenBerg et al.~(1983, VdB83). The range of predicted secondary star masses by these authors is rather large, ranging from 0.149\\,\\msun (CP90) and 0.174\\,\\msun (N84) to 0.212\\,\\msun (VdB83). The intersection points between both ($Q,i$)-relations shown in the upper panel have to be reflected at the corresponding lines in the second panel in order to read white-dwarf mass for the nominal orbital velocity. This yields a rather low white-dwarf mass, \\mwd $\\simeq 0.42 - 0.48$\\,\\msun. The high-mass limit is subject to the adopted mass-radius relation of the secondary star and the confidence level one would like to reach. For the reader's convenience we show the three by the different $M/R$-relations predicted orbital velocity amplitudes for the secondary star as a function of $Q$ in the lower panel of Fig.~\\ref{qirel}. For 1, 2, 3 $\\sigma$-errors of the radial velocity amplitude (335, 385, 435\\,\\kmps) and the mass-radius relation N84 one obtains $M_{\\rm wd, max} = 0.6, 0.80, 1.04$\\,\\msun, respectively. As an illustration we show with long dashes the corresponding mass function for the 2$\\sigma$-level in the upper panel of Fig.~\\ref{qirel}. The new NaI measurements suggest a white-dwarf mass not in excess of 1\\,\\msun\\ and a mass ratio not in excess of 6. Similarly the H$\\alpha$-line can be used to estimate the white dwarf mass. Its photometric and radial velocity variation suggest an origin on that hemisphere of the companion star which is illuminated by EUV-radiation from the accretion spot. The measured velocity amplitude has to be multiplied by a $Q$-dependent factor in order to scale the observed center-of-light velocity to the required center-of-mass velocity ($K_2$-correction). It is clear that the corresponding mass estimate must yield a higher value than that using the NaI lines, since both measured velocity amplitudes are more or less the same. A graph of the factor mentioned is shown in Schwope et al.~(1993), for $Q= 5.5$ its value is 1.35. The corresponding $(Q,i)$-relation is shown with a dash-dotted line in Fig.~\\ref{qirel} (upper panel) and suggests, as expected, a much higher mass of the white dwarf \\mwd $ \\simeq 0.95$\\,\\msun\\ ($Q \\simeq 5.5$) than obtained using the NaI line. The model for the $K_2$-correction assumes that reprocessed emission originates from the whole illuminated part of the secondaries Roche lobe and that its intensity at given velocity (surface element) is proportional to the locally incident radiation. The real situation might be different and the mass estimate therefore in error. We have found recently two pronounced examples for such a deviation in other AM Her systems. First, the H$\\alpha$ emission line radial velocity in MR Ser seen in a low state is the same as that of the NaI absorption line (as here in UZ For). This velocity amplitude is much higher than that of the narrow emission line from the secondary star seen in MR Ser's high state (Schwope et al.~1993). Second, Doppler tomography of the Balmer and HeII emission lines of HU Aqr observed in the high accretion state show that the Balmer emission is concentrated away from the inner Lagrangian point (Schwope et al.~1996). Both examples suggest that the reason for the negligible difference between the NaI and H$\\alpha$ velocity amplitudes in the present case of \\uz\\ is a similar deviation of the Balmer emission from the simple pattern and that, hence, the $K_2$-correction applied might be too large. A common solution at the $\\sim$2$\\sigma$-level for both our radial velocity measurements is $Q=4.4$, \\mwd $ = 0.75$, and $i = 80.7\\degr$, which is also well compatible with the photometric study by Bailey \\& Cropper (1991). Finally we comment on the possible origin of the higher radial velocities derived in Paper I. These data were obtained in a similar state of accretion and with comparable time and spectral resolution, but with a smaller telescope and with a high read-noise RCA CCD (as the data presented here). This allowed measurement of the NaI-lines in only a few spectra. It is possible that these measurements were misleading. The number of useful H$\\alpha$ measurements in both observations is the same, and the radial velocities are not too discrepant ($395\\pm67$\\,\\kmps at 90\\% confidence in Paper I, $308 \\pm 27$\\,\\kmps (1$\\sigma$) here). Since the H$\\alpha$-line measures the center of light on the illuminated hemisphere of the secondary with possible contaminations from the accretion stream a discrepant radial velocity amplitude would simply mean, that the center of light is shifted between the different observations. The degree of freedom for shifts of the NaI radial velocity, however, is smaller, although illumination might shift the center of light towards the non-illuminated hemisphere (e.g.~Davey \\& Smith 1992). Our main results are summarized: (1) The mass of the white dwarf in \\uz\\ is most likely not in excess of 1\\msun\\ ($i < 83\\degr, Q < 6$). (2) The mass-radius relation of low-mass main sequence stars introduces large uncertainties to mass estimates of white dwarfs in polars and needs proper calibration. (3) The ultimate location of the low-state Balmer emission on the illuminated secondary in \\uz\\ remains uncertain, hence the applicability of the $K_2$-correction scheme is questionable. One needs high signal-to-noise data with high phase resolution in order to calibrate empirically the $K_2$-correction. (4) The likely low mass of the white dwarf in UZ For together with its measured orbital period suggest that the system was born in or slightly above the period gap." + }, + "9701/astro-ph9701172_arXiv.txt": { + "abstract": "Gamma-ray bursters emit a small fraction of their flux in X rays, and because X-ray detectors are often very sensitive they may probe the gamma-ray burst universe more deeply than the current best gamma-ray instruments. On the reasonable assumptions that spectra of bursts observed by BATSE may be used to predict the X-ray fluxes of gamma-ray bursts, and that any corona of bursts around M31 is similar to the one around the Milky Way, we predict the rate at which the wide field cameras on board BeppoSAX should detect bursts from the Milky Way and M31. These rates are such that a one-month observation of M31 would have to either detect bursts from M31 or exclude most galactic models of gamma-ray bursts. (It is shown how the remainder can be dealt with.) Therefore such an observation would settle the long-standing dispute over their location. ", + "introduction": "\\label{intro} The results of the BATSE mission (see Fishman and Meegan 1995)\\nocite{fm:95}, combined with earlier data sets for bright bursts such as the one collected by PVO (Fenimore et~al.\\ 1993)\\nocite{fehkl:93} have shown that (1) gamma-ray burst positions are distributed uniformly and randomly on the sky (Briggs et~al.\\ 1996a,b)\\nocite{bppmf:96,bppmf2:96} and (2) the cumulative number as a function of peak flux, $N(>P_\\gamma)$, is consistent with a constant rate density of bursts within some volume around us, and a decreasing density outside that volume. This implies that we are at the centre of a gamma-ray burst universe of which we can see the edge and which looks the same in all directions. Most distance scales are therefore excluded. The first remaining one is the high-redshift universe, with the edge being caused either by cosmological volume effects near and beyond $z=1$ or by evolution of the density at moderate redshift (or both). The second one is an extended corona of our Galaxy, much bigger than the dark-matter halo and invented for the purpose of housing gamma-ray bursts (GRB). We are not strictly in its centre, but the average GRB distance can be made large enough that the anisotropy due to our offset from the centre is below the limit set by the BATSE data on burst positions. At the same time it can still be small enough that we need not see M31 (Briggs et~al.\\ 1996b). The aim of this paper is to demonstrate the capability of the Wide Field Cameras (WFC) on board BeppoSAX (launched in April 1996) to distinguish between these options by searching for the hypothetical corona of GRB around M31. We first discuss the Z-ray detectability of GRB (Sect.~\\ref{xgrb}) and our implementation of corona models (Sect.~\\ref{coro}). Our results are presented in Sect.~\\ref{resu} and compared to previous results in Sect.~\\ref{disc}. ", + "conclusions": "\\label{disc} Previous work on constraining coronal models of GRB has many similarities with our own calculations. Liang (1991)\\nocite{liang:91} found that ROSAT might detect some GRB in X rays assuming a now abandoned disc model for the distribution of GRB. Li, Fenimore, \\& Liang (1996)\\nocite{lfl:96} used a similar method to our own both for beamed and unbeamed models. Their calculations differ from ours mostly in that they use neither the spread in X-ray luminosity nor the X-ray excesses. Their hypothetical instrument had $S\\sub{min}=0.1\\photcms$. As we can see from \\Fig{fi:smin}, 30--90\\% of the bursts that the WFC can see on-axis are below this limit, so it is no surprise that we find more optimistic prospects for detecting GRB in M31. Harrison and Thorsett (1996)\\nocite{ht:96} considered a variety of real instruments, calculating the detectable rate in much the same way as we did (including the spectral variability using the same set of spectra from Band et~al.\\ 1993). They conclude that only a novel instrument sensitive to photons in the 10--200\\,keV range and with a field of view of 18$^\\circ$ would be capable of detecting M31 in one year. While they did not include the possibility of X-ray excesses, they would without doubt have realised the potential of the SAX WFC if they had included them in their work. In summary, we have shown that the hitherto neglected spread in X-ray to gamma-ray luminosity ratios of gamma-ray bursts substantially increases the prospects for deciding the gamma-ray burst distance scale. The case is further improved greatly by the recent discovery that a substantial fraction of gamma-ray bursts have X-ray excesses (Preece et~al.\\ 1996). A one-month observation of M31 with an existing instrument, the SAX WFC, will be decisive for establishing whether or not the Andromeda Nebula harbours a population of bursters, unless bursters only emit radiation in fairly narrow cones along their direction of motion. In that case, a dedicated, cheap mission similar to the SAX WFC should resolve the issue in about one year of observing time. Observing proposals to do the experiment in WFC secondary (i.e.\\ unguaranteed) time have been accepted, so the gamma-ray burst distance scale may not remain uncertain much longer." + }, + "9701/astro-ph9701034_arXiv.txt": { + "abstract": "In hierarchical clustering, galaxy clusters accrete mass through the aggregation of smaller systems. Thus, the velocity field of the infall regions of clusters contains significant random motion superimposed on radial infall. Because the purely spherical infall model does not predict the amplitude of the velocity field correctly, methods estimating the cosmological density parameter $\\Omega_0$ based on this model yield unreliable biased results. In fact, the amplitude of the velocity field depends on local dynamics and only very weakly on the global properties of the universe. We use $N$-body simulations of flat and open universes to show that the amplitude of the velocity field of the infall regions of dark matter halos is a direct measure of the escape velocity within these regions. We can use this amplitude to estimate the mass of dark matter halos within a few megaparsecs from the halo center. In this region dynamical equilibrium assumptions {\\it do not} hold. The method yields a mass estimate with better than $30\\%$ accuracy. If galaxies trace the velocity field of the infall regions of clusters reliably, this method provides a straightforward way to estimate the amount of mass surrounding rich galaxy clusters from redshift data alone. ", + "introduction": "The linear theory of density perturbations shows that a spherically symmetric mass concentration in an expanding universe induces a radial peculiar velocity field in the surrounding region \\begin{equation} {{v_{\\rm pec}(r)}\\over {H_0r}} = -{{1}\\over {3}}\\Omega_0^{0.6} \\delta(r) \\label{1.1} \\end{equation} where $\\delta(r)$ is the average spherical mass overdensity within the radius $r$, and $H_0$ and $\\Omega_0$ are the Hubble constant\\footnote{We use $h=0.5$ throughout, where $H_0=100h$ km s$^{-1}$ Mpc$^{-1}$.} and the cosmological density parameter at the present time, respectively. We can estimate the galaxy number overdensity $\\delta_g=b\\delta$ where $b$ is the bias parameter. Thus, if we can measure $v_{\\rm pec}$ we can estimate $\\beta=\\Omega_0^{0.6}/b$. Several authors have applied this method to the Local Supercluster and to the Virgo cluster with mixed results (see e.g. the review by \\cite{Davis83}; see also \\cite{Strauss95}). Along with the uncertainties in the determination of the galaxy number overdensity and the peculiar velocity themselves, external tidal shear can strongly affect the velocity field (e.g. \\cite{Hoff86}; \\cite{Eisen95}; \\cite{Bond96}); thus we cannot obtain reliable results unless we sample the velocity field within the full three-dimensional region around the cluster (\\cite{Vill86}). In redshift space, spherical infall confines galaxies around clusters within caustics, surfaces with a characteristic ``trumpet'' shape (\\cite{Kaiser87}). For these non-linear regions, we must replace the linear regime equation (\\ref{1.1}) with the exact solution of the equation of motion of shells in an overdense region within the expanding universe (e.g. \\cite{Silk77}). However, the dependence of the peculiar velocity on $\\beta$ and $\\delta_g$ is still approximately separable (\\cite{Reg89}). Thus, provided we can determine the location of the caustics, we can still estimate $\\beta$ through the measurement of $\\delta_g$ for a rich galaxy cluster (\\cite{Reg89}). Recently, \\cite{Reg96} suggested the application of gravitational lensing (e.g. \\cite{Tyson90}; \\cite{Kais93}; \\cite{Kais95}; \\cite{Bonn95}) to determine the mass overdensity $\\delta$ and therefore the unbiased value of $\\Omega_0$. Standard inflationary cosmologies (e.g. \\cite{Peacock96}) predict that the primordial density field is a Gaussian random field. These initial conditions can lead to either top-down or bottom-up scenarios for the formation of cosmic structures. The hierarchical clustering in the bottom-up scenarios, with large systems forming by aggregation of smaller ones, currently represents the most successful framework of structure formation theories as it seems to be able to reproduce many, although not all, properties of the real universe. The top-down scenarios are far less succesfull (see e.g. \\cite{Pad93} for a general discussion). The general formation process of clusters in hierarchical scenarios differs substantially from the one described by the idealized spherical model; here clusters form through the infall of smooth spherical shells onto an initial density peak (\\cite{Gunn72}). Despite its idealized nature, the spherical model appears to represent a reasonable description of the collapse of high peaks in a random Gaussian density field (\\cite{Bern94}), and predicts density and velocity profiles of final systems in reasonably good agreement with $N$-body simulations of dark matter halo formation when the power spectrum has an effective spectral index $n\\ge -1$ (\\cite{Zar96}). These authors suggest that particle ranks in binding energy are conserved during the formation process. This conservation is responsible for the agreement (see also \\cite{Hoffman88}; \\cite{Quinn88}; \\cite{Zar93}). In fact, for spectral indices $n\\le -1$ the formation process is more violent (\\cite{LyndenBell67}), energy ranks are not conserved, and the agreement breaks down. Thus, the agreement in energy space between the spherical infall model and hierarchical clustering for a limited range of spectral index, $n$, allows correct predictions about the {\\it final state} of the halo. However, the spherical infall model does not describe the {\\it evolution} of the outer regions of systems in configuration space. $N$-body simulations of flat universes (\\cite{Haar92}; \\cite{HaarWeyg93}) show that the velocity field of cluster surroundings is very different from the predictions of spherical infall because (1) recent mergers increase the particle kinetic energy, and (2) the presence of substructures makes the velocity profile irregular. In other words, random motions obscure the infall information. If the same problem arises in real galaxy clusters, estimates of $\\beta$ (or $\\Omega_0$) based on spherical infall are systematic overestimates. Here, we suggest a unifying explanation for the amplitude of the caustics. We use $N$-body simulations of flat and open universes to show that the escape velocity around dark matter halos determines the amplitude of the caustics. Thus, the local dynamics of the halo, not the global properties of the universe, dominates the amplitude of the velocity field around clusters. We also show that under reasonable hypotheses about the density field outside the virialized region, the velocity field around a cluster provides an estimate of the mass enclosed within a few megaparsecs of the cluster center. This method is particularly interesting because, in this region, the equilibrium assumptions underlying usual mass estimation methods do not hold. In Sect. 2 we review the main equations of the spherical infall model. In Sect. 3 we derive the alternative expression for the amplitude of the caustics and we outline the mass estimation method. In Sect. 4 we compare both the expression for the amplitude of the caustics and the spherical infall model with the results of $N$-body simulations. We also use these $N$-body simulations to test our new mass estimation method. ", + "conclusions": "In redshift space, galaxies around clusters should appear within regions with a characteristic trumpet shape (\\cite{Kaiser87}). Reg\\\"os \\& Geller (1989) % suggested applying the spherical infall model to these caustics to constrain the density of the universe. We show that this method generally overestimates the actual density parameter $\\Omega_0$ because random motions increase the amplitude of the caustics (see also \\cite{Lilje91}; \\cite{Zaritsk92}). We define the amplitude of the velocity field as half of the difference between the maximum and minimum line-of-sight velocity at projected distance $r_\\perp$ from the cluster center. We use $N$-body simulations to show that the escape velocity (eqs. [\\ref{3.2}], and [\\ref{4.4}]) describes this amplitude well, with $r_\\perp$ ranging over two orders of magnitude, from the central region of the halo to its infall regions, where particle orbits are mainly radial. Van Haarlem (1992) % first noted that the spherical infall method overestimates $\\Omega_0$ and pointed out that mergers and substructures within the infall regions are responsible for the disagreement. Here we suggest a unifying explanation. We show that our interpretation of the amplitude of the velocity field within halo infall regions can be applied to estimate the interior mass of halos up to a few virial radii $r_\\delta$ from the halo center, where the usual equilibrium assumptions do not hold. This estimation technique works because the local dynamics depends more strongly on the mass of the halo than on the global properties of the universe. This deduction agrees with the recent suggestion by White (1996) % and \\cite{Navarro96} that halos have a universal density profile, basically independent of the cosmology, with a characteristic density depending on the formation time of the halo (see also our eq. [\\ref{4.2}]) which ultimately depends on the halo mass. In the literature, the dependence of the density profile on the halo mass has been neglected, with attention focused on the dependence on $\\Omega_0$ and on the power spectrum $P(k)$. The limited overdensity range examined in these numerical studies explains why $N$-body simulations appear to show a dependence of the halo density profile on the underlying cosmology (e.g. \\cite{Crone94}). If we know the mass $M( = N_{HI}/(1.823{\\times}10^{18}\\int \\tau_{21}(v) dv)$ (cf. Dickey \\& Brinks \\markcite{D1} 1988), yields an estimate for $ \\approx 1380$~K. This temperature is the column-density-weighted, harmonic mean of the spin temperature along the sight line (Kulkarni \\& Heiles \\markcite{K1} 1987, Dickey \\& Lockman \\markcite{D2} 1990). $ \\approx 1380$~K is substantially higher than is observed along lines of sight through the Milky Way Galaxy or M31 (Dickey \\& Brinks \\markcite{D1} 1988), where $ \\sim$ 200 to 300~K are typical for column densities $N_{HI}\\approx 10^{21}$cm$^{-2}$. The estimate for $T_s$ in the $z=3.4$ absorber does have uncertainty in addition to those attached to $\\tau_{21}$ and $\\Delta V$; we estimate the uncertainty in $N_{HI}$ to be ${\\sim}$30\\%, which is not enough to alter the conclusion that $$ computed in this way is high. We have assumed that all the neutral gas that contributes to the damped Lyman~$\\alpha$ profile is confined to the Gaussian feature that we observed in the 21--cm line. A sharply contrasting possibility is that the high redshift system has substantial quantities of warm, low optical depth gas that is sufficiently spread in velocity to escape inclusion in the 21--cm profile. Then, the inferred $T_s$ for the 21--cm absorber becomes ${\\simeq}T_sf_{21}$, where $f_{21}$ is the fraction of $N_{HI}$ that is contained in the 23 km~s$^{-1}$ profile (velocity dispersion $\\sigma \\approx 10$ km~s$^{-1}$). Thus, in addition to the narrow 21--cm absorption component, a complete description of the kinematics of the system must include the broader, low column density component whose cloudlets are optically thick in the resonance lines of common metals but optically thin in the 21--cm line. In other systems at lower $z$, this component has been tentatively identified with halo gas that encases the dynamically cold, 21--cm absorbing disk (Wolfe \\& Wills \\markcite{Wo2} 1977, Briggs \\& Wolfe \\markcite{Br1} 1983, Briggs et al \\markcite{Br2} 1985, Lanzetta \\& Bowen \\markcite{L1} 1992). For these cases, high-resolution optical spectroscopy has led to the general conclusion that most of the neutral column density in the damped Lyman~$\\alpha$ systems is concentrated in the cold, low-dispersion component (c.f. Wolfe et al \\markcite{Wo5} 1994), and only a tiny fraction of the neutral gas is contained in the broad, turbulent (or ``halo'') component, which dominates in providing the equivalent widths measured for the metal lines. Thus, $f_{21}$ is likely to be close to unity, and $T_s \\approx 1380$~K as indicated above. For this system in absorption against MG0201+113, the effective velocity spread of the broad component can be estimated from the equivalent widths of the metal lines observed in the low resolution (10\\AA) spectra of White, Kinney and Becker \\markcite{W2} (1993). We performed a curve of growth analysis using the line depths and limits to depth for several transitions of SiII ($\\lambda$1260, $\\lambda$1526, $\\lambda$1808), FeII ($\\lambda$1144, $\\lambda$1608), CII ($\\lambda$1334), and AlII ($\\lambda$1670). Although there is considerable uncertainty in this analysis, it serves to illustrate the constraints on the parameters that need to be specified to better determine the spin temperature. A range of two component models is illustrated in Figure 4: the cold cloud has $\\sigma_{21} = 10$ km~s$^{-1}$ and $N_{21}=f_{21}N_{HI}$, and the ``halo'' component is modeled as a single broad dispersion cloud with $\\sigma_{ha}$ and neutral hydrogen column density $N_{ha}=f_{ha}N_{HI}=(1-f_{21})N_{HI}$. Due to the scarcity of information, a single measure for the metal abundances relative to solar abundances ``$X$'' is assumed for all ions in both components; for example, the total column density of CII in the system is specified by $N_{CII}= 10^XN_{HI}{\\times}[n_C/n_H]_{\\odot}$, where $[n_C/n_H]_{\\odot}$ is the solar abundance of carbon relative to hydrogen as tabulated by Spitzer \\markcite{S2} (1978). This approach also assumes that the carbon associated with the $N_{HI}$ in this layer is predominantly singly-ionized, in keeping with neutral clouds along lines of sight through the Milky Way Galaxy. The upper panels of Figure 4 illustrate models in which the 21--cm absorber dominates the neutral column density; the lack of detection of the SiII$\\lambda$1808 line enforces a low metal-abundance estimate, and the stronger detections of other metal lines indicates that there is additional, turbulent gas. The data permit the velocity spread of the turbulent component to be very broad, although the CII line would be resolved by the instrumental resolution if the effective $\\sigma$ were much over 200~km~s$^{-1}$. The low resolution spectroscopy is compatible with a wide range of models. For example, the metal line strengths are also compatible with much larger fractions of gas lying in the broad component, as is shown in the lower panels of Figure 4; this range of models must have much lower metal abundances. { \\plotone{Briggs.fig4.ps} \\figcaption[Briggs.fig4.ps]{Curve of Growth (COG) display of the range of two-component kinematical models that simultaneously fit both the 21--cm line profile and the ultraviolet lines. The {\\it top panels} illustrate the range of models with the highest metal abundances (relative to solar abundances) and lowest fraction of neutral gas in the broad velocity dispersion, ``halo'' component. The {\\it lower panels} explore the range of models with the large fractions of gas in the ``halo'' component, yielding the lowest values of metal abundances. The dashed COG represents the $\\sigma_{21} = 10$ km~s$^{-1}$ 21--cm absorption component. The solid line is the two component model. The damping portion of the COG is drawn for Lyman~$\\alpha$. The abscissa has units of cm$^{-1}$. } } The constraints on the relative column densities in the turbulent and cold components can be translated into measurements of the spin temperatures for both. These spin temperatures and the other model parameters from the COG analysis are summarized in Figure 5. The metal abundances from the models range from $10^{-1.6}$ to $10^{-2.7}$ below solar. The lower limit to the spin temperature of the turbulent halo component results from inspecting the spectra in Figures 2 and 3 and noting that an absorption-line three times broader than the detected line would probably escape detection if its depth were less than $\\sim$5~mJy. Although $f_{21}$ is expected to be close to unity, the certainty of the interpretation would benefit from new, high-resolution optical spectroscopy that can demonstrate that the bulk of the gas is confined to the 21--cm line velocity width. { \\plotone{Briggs.fig5.ps} \\figcaption[Briggs.fig5.ps]{Constraints on spin temperature as a function of the fraction of the neutral gas contained in the 21--cm component, $f_{21}$. The fraction of gas contained in the broad velocity dispersion ``halo'' component, $f_{halo} = 1 - f_{21}$, is plotted across the top horizontal axis. {\\it Upper panels} indicate the range of effective velocity dispersions, $\\sigma$, and lower limit to spin temperature, $T_{spin}$, for the ``halo'' component, determined from the COG analysis of Figure 4 and the detection limits in the 21--cm line. The {\\it lower-central panel} plots the logarithm of the metal abundance relative to solar, assuming that it is the same in both components. The {\\it lower panel} plots the spin temperature of the 21--cm line component as a function of $f_{21}$. } } Further uncertainty in determination of the spin temperature results from possible differences in the spatial extent of the absorber and the background radio source, which may be substantially larger than the source of optical continuum causing a significant portion of the radio flux to leak around the thick absorbing cloud that creates the damped Lyman~$\\alpha$ line. The continuum map at 18cm wavelength (Stanghellini et al \\markcite{S3} 1990) hints that this might be the case by showing evidence for a weak component extending to ${\\sim}4''$ from the nucleus, although the extended emission accounts for only a tiny fraction of the total radio flux. While MG0201+113 falls in the GHz--Peaked-Spectrum class of radio source, the integral flux density spectrum for the source shows a flattening at meter wavelengths (O'Dea et al 1990), indicating that at the frequency of the redshifted 21--cm absorption line, the source is dominated by a more extended component than the compact optically thick core component that gives rise to the GHz peak in the continuum spectrum (de Bruyn et al 1996). High resolution 90cm VLBI observations are capable of defining the continuum source structure, which ultimately will clarify the interpretation. If the resulting $$ remains high, then the neutral medium in this absorber is substantially hotter than the neutral ISM of the Milky Way or M31. This may indicate that a larger fraction of the total HI is in the Warm Neutral phase with $T \\approx 8000$~K (Kulkarni \\& Heiles \\markcite{K1} 1987), or alternatively, the gas in the Cold phase that we detect in the 21--cm line may actually be hotter due to the inability of the gas to cool itself when the metal abundances are low. The new Arecibo HI spectroscopy also places limits on the emission from large, neutral gas masses that might be associated with the damped Lyman~$\\alpha$ site. The sensitivity is a strong function of the velocity width of the signal (cf. Wieringa et al. \\markcite{W1} 1992), which would be expected to range from a few hundred km~s$^{-1}$ for pancakes (Zel'dovich 1970) to more than 1000 km~s$^{-1}$ for analogs to present day galaxy clusters. Two examples of Gaussian profiles for $\\Delta V = 200$ km~s$^{-1}$, $M_{HI} = 10^{13}\\msolar$ and $\\Delta V = 1000$ km~s$^{-1}$, $M_{HI} = 5{\\times}10^{13}\\msolar$ ($H_o$ = 100 km~s$^{-1}$~Mpc$^{-1}$ and $\\Omega_o = 1$) are plotted in Figure 1d for comparison with the data of Figure 1c, after averaging to 100 kHz wide spectral bins. Only a single linear baseline has been subtracted from the entire 2.5 MHz band. The choice of the narrow 2.5 MHz band for this experiment limits would prevent the detection of $\\Delta V$ $\\lower .7ex\\hbox{$>$}\\atop \\raise .2ex\\hbox{$\\sim$}$ 1200 km~s$^{-1}$. Clearly, systematic effects, such as a weak, characteristic ``standing wave'' and low level RFI, dominate in limiting the sensitivity to broad signals. On the other hand, signals as large as these would be noticeable. The limits are still more than a factor of 10 higher than Wolfe's \\markcite{Wo4} (1993) prediction for the amount of HI associated with damped Lyman~$\\alpha$ sites." + }, + "9701/astro-ph9701129_arXiv.txt": { + "abstract": " ", + "introduction": "According to the traditional view galaxies are clustered on small scales, and on large scales the distribution of galaxies and clusters of galaxies is random (i.e. scale-free). Such a distribution is expected for the currently popular structure formation scenario based on the evolving dynamics of the dark matter during the history of the Universe. Quantitatively this behaviour is described by the correlation function of galaxies and clusters of galaxies which has a high peak at small separations but approaches zero at about 30 \\Mpc\\ for galaxies and at about 70 \\Mpc\\ for clusters of galaxies (in this paper we express the Hubble constant as $H_0=h~100~{\\rm km~s^{-1}~Mpc}^{-1}$). Initially these scales were thought to be the scales of the transition to homogeneity. However, in the distribution of galaxies voids of diameter up to 50 \\Mpc\\ were discovered, voids in the distribution of clusters of galaxies are even larger with diameters of 100 \\Mpc\\ and more. Thus it was thought that the homogeneity begins on supercluster scale. Recent studies have shown that even on supercluster scale the distribution of galaxies and clusters of galaxies may have some regularity. First clear indication for the presence of periodicity in the distribution of high-density regions in the distribution of galaxies came from the deep pencil-beam survey of redshifts of galaxies by Broadhurst \\etal (1990). High-density regions in this survey form a regular pattern with a period 128 \\Mpc. Nearest peaks in the distribution of high-density regions in this survey coincide in position and redshift with superclusters of galaxies (Bahcall 1991), thus one may think that the distribution of superclusters may have some regularity. However, as no clear periodicity was found in other directions, this result was explained as a statistical anomaly (Kaiser and Peacock 1991). Independent studies have shown that approximately on the same scale the cluster correlation function has a weak secondary maximum (Kopylov \\etal (1988), Mo \\etal (1992), Einasto and Gramann (1993), Fetisova \\etal (1993)). This scale has been detected in the three-dimensional supercluster-void network by Einasto \\etal (1994), Einasto (1995a, 1995b). The analysis of the Las Campanas Redshift Survey has shown that the distribution of sheet-like and filamentary structures has a preferred scale about 100~\\Mpc\\ (Doroshkevich \\etal 1996), also the two-dimensional power spectrum of galaxies of this survey has a peak on the same scale (Landy \\etal 1996). These results were based on small number of objects or (with the exception of the study by Einasto \\etal 1994) on two-dimensional data, thus further studies are needed. To investigate the distribution of matter we have used a new redshift compilation of rich clusters of galaxies by Andernach, Tago and Stengler-Larrea (see 1995). This is the largest and deepest three-dimensional survey available presently. Using this dataset we have compiled a new catalogue of superclusters of galaxies (Einasto \\etal 1997a), calculated the nearest neighbour and void diameter distribution (Einasto \\etal 1997a), and determined the cluster correlation function and power spectrum (Einasto \\etal 1997b, 1997c, 1997d). Here I give a short summary of principal results of these studies. ", + "conclusions": "Our study of the distribution of clusters of galaxies has lead us to the following main conclusions. \\begin{itemize} \\item{} The distribution of high-density regions in the Universe (rich superclusters) is more regular than expected previously. Superclusters and voids form a cellular lattice or network with step size $120 \\pm 20$~\\Mpc. The location of cells is rather regular. \\item{} The correlation function of clusters of galaxies has an oscillatory behaviour with regularly spaced secondary maxima and minima. The period of oscillations, 120 \\Mpc, is equal to the scale of the supercluster-void network. The power spectrum of the cluster correlation function has a sharp peak on the respective wavelength. \\item{} Clusters of galaxies in CDM-type models of structure formation are located less regularly than real clusters. \\item{} If the distribution of clusters of galaxies reflects the distribution of all matter then presently popular structure formation theories need revision. \\end{itemize} ACKNOWLEDGEMENTS. I thank my colleagues H. Andernach, M. Einasto, P. Frisch, S. Gottl\\\"ober, V. M\\\"uller, V. Saar, A. A. Starobinsky, and E. Tago for fruitful collaboration and for the permission to use our results prior to publication. This work was partly supported by the grants from the International Science Foundation, Estonian Science Foundation, German Science Foundation and Russian Federation." + }, + "9701/astro-ph9701153_arXiv.txt": { + "abstract": "From deep color--magnitude arrays made from $V$ and $I$ images taken with {\\it Hubble Space Telescope\\/}'s WFPC2 camera we have determined luminosity functions (LFs) down to a level that corresponds to $\\sim$0.13 \\msun, for the low-metal-abundance globular clusters M15, M30, M92, and NGC 6397. Because of the similarity of the metallicities of these clusters, differences in their luminosity functions directly trace differences in their mass functions. The LFs of M15, M30, and M92 agree closely over the entire observed range, whereas that of NGC 6397 drops away sharply at the faintest magnitudes. We suggest that the deficiency of low-mass stars in NGC 6397 is due to tidal shocks, to ejection through internal relaxation, or to a combination of the two. With the presently available mass--luminosity relations, we find that even in M15, M30, and M92 the mass functions probably do not rise so fast as to make the low-mass end dominant. ", + "introduction": "The {\\it Hubble Space Telescope} (\\hst) now makes it possible to derive color--magnitude diagrams (CMDs) of globular clusters (GCs) that extend several magnitudes fainter than the CMDs found for the same clusters from the ground. Among the important by-products of these CMDs are luminosity functions (LFs), and from them mass functions (MFs), that can extend over nearly the entire length of the lower main sequence. Insight into the formation and dynamical evolution of globulars can be gained by using these LFs and MFs to compare the stellar content of different clusters. It has often been suggested that the mass function of a cluster depends on its position in the Galaxy, its metallicity, and its dynamical history (see, e.g., McClure \\etal\\ v1986, Aguilar, Hut, \\& Ostriker 1989, Djorgovski, Piotto, \\& Capaccioli 1993); and the relationships are complicated by the fact that the latter two factors correlate with the first. In any case, the main-sequence mass functions of clusters are one of our important clues relating to their origin. Furthermore, globular clusters have quite often been used as tracers and indicators of the halo population of which they are the most conspicuous part. In particular, it has been suggested that there is a steep rise at the low-mass end of cluster MFs (Richer \\etal\\ 1991), and that this might be indicative of a considerable contribution by low-mass stars to the little-known total mass that governs the poorly understood flat rotation curve of the Galaxy (Richer \\& Fahlman 1992). In the case of the Galaxy, the steep MF slope that Richer and Fahlman (1992) found at the faint end has been contradicted by better data by Reid \\etal\\ (1996); and we shall show here that their assertion about steep MFs for low-mass stars in globular clusters also fails to be borne out by improved data that we will present below. Of equal importance will be the demonstration that with one exception that is probably understandable, our sample of clusters shows a remarkable uniformity of mass functions. Progress on these issues requires accurate photometry of low-mass main-sequence stars. Ground-based studies have so far been limited largely to the mass range above 0.4\\msun; deeper studies have been pursued only with difficulty and uncertainty. With the advent of \\hst, however, we can largely circumvent the crowding problems that plague ground-based observations, reaching several magnitudes deeper, down to masses close to the bottom end of the hydrogen-burning main sequence. Here we present the first deep \\hst\\ LFs of M30 (NGC 7099) and M92 (NGC 6341), and independent measurements of the LFs of NGC 6397 and M15 (NGC 7078). The CMD and the LF of NGC 6397 have already been presented in Cool, Piotto, \\& King (1996, CPK). We compare our results for the latter two clusters with \\hst\\ LFs of Paresce, De Marchi, \\& Romaniello (1995) and De Marchi \\& Paresce (1995), and then compare the four cluster LFs with each other. The results of our comparison differ from those reached by De Marchi \\&\\ Paresce (1995), in that we find a significant difference between the LFs of NGC 6397 and M15. Two features of this set of four clusters make for a particularly useful comparison. First, three of them have collapsed cores and similar surface-brightness profiles, which make the conversion from the local, observed LFs to global LFs straightforward. M92 is a high-concentration cluster; in this case also the observed LF does not differ appreciably from the global one. Second, all four clusters have comparable metallicities (see Table 1), so that similarities or differences in their LFs will directly reflect similarities or differences in their MFs. A description of the observations and of the methods of analysis are given in Section 2. We present the cluster LFs and MFs in turn in Sections 3 and 4, exploring in the latter the role of the mass--luminosity relation. The findings are summarized and discussed in Section 5. ", + "conclusions": "" + }, + "9701/astro-ph9701015_arXiv.txt": { + "abstract": "Using imaging \\fp\\ data, we study the star-forming properties and kinematics of the nearby barred spiral galaxy NGC 2442. The \\ha emission is very localized along the strong spiral arms of the galaxy, and shows a marked asymmetry between the sharp, well-defined northern tidal arm and the weaker southern arm. The \\ha velocity field appears highly distorted, with a rapidly rotating nuclear component. We find evidence for strong non-circular motions along the northern arm, coincident with the pronounced dust lane and regions of intense star formation. The strong asymmetries, disturbed velocity field, and presence of a perturbed companion suggest that we are witnessing a strong kinematic response to a close interaction, which has redistributed the star formation activity throughout the disk of NGC 2442. Dynamical modeling of the NGC 2442 system supports this hypothesis, and suggests that the regions of strongest star formation are coincident with strong shocks occurring along the tidally perturbed northern arm. Despite this strong redistribution of the gas on small scales, this galaxy does not show a significant departure from the Tully-Fisher relation, nor does it appear to be experiencing any strong starburst. Moreover, our models predict that in a few x 10$^{8}$ years, NGC 2442 will have recovered from this first tidal encounter and will experience another passage -- and ultimately a merger -- in a few Gyr. This merger may provoke stronger, permanent changes in the structural properties of the galaxy, depending on the detailed response of the disk. Given the environment of many disk galaxies, this tidal encounter cycle seems likely to be a normal phase of disk galaxy evolution. ", + "introduction": "There is compelling morphological and photometric evidence that enhanced star-forming activity in disk galaxies is a direct result of tidal interactions between galaxies. However, the physics of this strong causal relationship remain poorly understood. Although the strongest starbursts are generally found in the nuclei of interacting galaxies, some systems only show star formation in their outer disks, while others show no sign of elevated star formation rates at all. This variety suggests the local physics of gas dynamical response to perturbations varies considerably from interaction to interaction. While models of interaction-induced star formation can successfully explain strong central starbursts through bar-induced inflows (\\eg Noguchi 1988; Mihos, Richstone, \\& Bothun 1992; Mihos \\& Hernquist 1994, 1996), they have more difficulty exciting strong disk star formation. Moreover, induced star formation in interacting disks is often asymmetric and patchy in nature, suggesting that the {\\it local} dynamical conditions in the galaxy are an important piece in the puzzle. To understand in detail how interactions influence star forming activity, we must simultaneously examine both the kinematics and the star forming properties of individual interacting systems. The southern galaxy NGC 2442, shown in Figure 1, is a nearby (D=16 Mpc for $H_0=75$ km/s/Mpc) peculiar barred spiral (SAB(s)bc, de Vaucouleurs \\etal 1991 [RC3]) showing many signs indicative of an interaction. Along with the short central bar, ovally distorted body, and isophotal twists, the disturbed, asymmetric, and extended arms of the galaxy argue for a collisional origin. Table 1 summarizes the environment of NGC 2442. The galaxy has two nearby companions and is possibly a member of a loose group (Tully 1988). Because of its proximity, NGC 2442 offers a rare opportunity to investigate the kinematic and star forming properties of an interacting galaxy shortly after the initial collision. At 16 Mpc, 1\\arcsec\\ subtends 77 pc, affording a very detailed view into the response of a galaxy to a collisional encounter. While a complete understanding of the dynamics of interacting galaxies is impossible when only morphological data exist, the combination of morphology and kinematic information can break the ambiguities arising from projection effects and yield a more accurate description of the dynamical state. Such an approach has been applied previously using kinematic information from slit spectra or HI data (\\eg Borne 1988; Borne \\etal 1988; Balcells \\etal 1989; Stanford \\& Balcells 1991; Hibbard \\& Mihos 1995). However, slit spectra do not provide full two-dimensional mapping of the ionized gas, while HI observations suffer from relatively low spatial resolution. The use of an imaging \\fp\\ spectrograph overcomes both these obstacles, providing two dimensional intensity and velocity maps of the ionized gas at $\\sim$ 1\\arcsec\\ resolution (\\eg Tilanus \\& Allen 1991; Vogel \\etal 1993, Canzian \\etal 1993; Mihos, Bothun, \\& Richstone 1993). Using the morphology and velocity field in conjunction with numerical modeling, it is possible to reconstruct the dynamical history of interacting galaxies (\\eg Mihos \\etal 1993) and address questions pertaining to the physical triggering mechanism for star-forming activity. Numerical models of colliding galaxies demonstrate clearly that strong interactions drive rapid dissipation and inflow in the gaseous components of galaxies (\\eg Barnes \\& Hernquist 1991, 1996; Mihos, Richstone, \\& Bothun 1992; Mihos \\& Hernquist 1994, 1996). This inflow is associated with the tidally-induced arms and bars in the galaxies, as gas is compressed along these features and driven inwards. In many galaxies, the radial motion along spiral features easily detectable: $\\sim$ 30 \\kms\\ for M81 (Visser 1980), $\\sim$ 50 -- 70 \\kms\\ for UGC 2885 (Canzian \\etal 1993), $\\sim$ 60 -- 90 \\kms\\ for M51 (Vogel \\etal 1988, Tilanus \\& Allen 1991). The fact that M51, closely interacting with NGC 5195, shows the highest radial gas motions lends credence to the theoretical models of strong inflows in galaxy interactions and suggests such motions should be easily detectable in other interacting galaxies given the resolution of \\fp\\ data. The velocity field of the ionized gas in and around tidal features should therefore yield a great deal of information about the strength of the inflow and the physical conditions in the gas. In this paper, we present imaging \\fp\\ data on NGC 2442 in order to explore the dynamical link between interactions, gaseous inflow, and star-forming activity. We find evidence for strong shocks and inflow along the galaxy's northern arm, accompanied by intense, localized star formation. Very little star formation is occurring in the southern tidal debris or along the bar itself, and in the inner kpc we detect a small, rapidly rotating ring of ionized gas. Based on these observations, we suggest a scenario in which NGC 2442 has recently (\\ie within a few hundred Myr) experienced a relatively strong interaction with \\am, driving the formation of the tidal arms. Dynamical modeling supports this scenario, and identifies the strong star forming sites along the arms as regions experiencing strong shocks and dissipation. The asymmetry observed in the star forming morphology of the galaxy seems to be related to the local dynamical conditions in the evolving tidal debris. ", + "conclusions": "Using \\fp\\ mapping of the morphology and kinematics of NGC 2442 in conjunction with a representative numerical model of the system, we have explored the triggering of star forming activity by a close interaction. Our results indicate that in the recent past, $\\sim$ 150 -- 250 Myr ago, NGC 2442 and \\am\\ experienced a close encounter, resulting in the tidal distortions currently observed in both galaxies. The distorted isophotes and velocity patterns throughout the disk of NGC 2442 attest to the damage done by even this relatively high mass ratio encounter. Material on the side of the disk closest to the companion experienced significant tidal shear, shredding much of the coherent tidal structure. Conversely, material on the far side of the disk was subject to a more coherent tidal compression, resulting in the formation of the dramatic northern tidal arm. As gas returns to the disk along this arm, it is being shocked and recompressed, forming a ridge of intense star formation and radial inflow along the NGC 2442's northern edge. The asymmetric pattern of star formation in NGC 2442 is a reflection of the differing physical conditions in the gas on either side of the galaxy. On the northern side of the galaxy, the gas is undergoing strong shocking and dissipation along the tidally induced arm, resulting in the vigorous, localized star formation. The extended tidal arm on the south side of the galaxy is much more diffuse, and the gas is not as strongly compressed along this feature. This kinematic structure is a reflection of the formation history of the tidal arms. Material close to the companion at perigalacticon experiences a very strong tidal shear, dragging material out into a transient tidal bridge which rapidly disperses (\\eg Toomre \\& Toomre 1972; Mihos \\etal 1992, 1993). In contrast, material on the opposite side of the disk feels a gentler -- but still effective -- tidal field, resulting in a more coherent tidal tail (or, in a high mass ratio encounter such as this, tidal arm). In this scenario, the south side of NGC 2442's disk was the closest point to the passing companion; in the subsequent time since the encounter, the outer disk has rotated significantly, now presenting the NE side closest to \\am in projection. The subsequent evolution of the NGC 2442 system is unclear. It is likely that the interaction will have extracted sufficient energy to bind the galaxies, and in our model the galaxies will merge in $\\sim$ 3 Gyr. In principle, this ultimate merger may be violent enough to lead to a transformation of NGC 2442 from Sbc to Sa or even S0, depending on the detailed structure of the merging galaxies. In the meantime, NGC 2442 will have many rotation periods to resettle into a relatively normal configuration. While the local gas dynamical response to these first encounters can be dramatic, the collisions don't seem to significantly alter the basic structure and global properties of large disk galaxies, until the orbit decays sufficiently to lead to a merger. We are most probably catching NGC 2442 in a rather transient but common phase of disk galaxy evolution -- recovering from a recent collision, and awaiting another yet to come." + }, + "9701/astro-ph9701147_arXiv.txt": { + "abstract": "We measure the topology (genus) of isodensity contour surfaces in volume limited subsets of the 1.2 Jy IRAS redshift survey, for smoothing scales $\\lambda=4\\hmpc$, $7\\hmpc$, and $12\\hmpc$. At $12\\hmpc$, the observed genus curve has a symmetric form similar to that predicted for a Gaussian random field. At the shorter smoothing lengths, the observed genus curve shows a modest shift in the direction of an isolated cluster or ``meatball'' topology. We use mock catalogs drawn from cosmological N-body simulations to investigate the systematic biases that affect topology measurements in samples of this size and to determine the full covariance matrix of the expected random errors. We incorporate the error correlations into our evaluations of theoretical models, obtaining both frequentist assessments of absolute goodness-of-fit and Bayesian assessments of models' relative likelihoods. We compare the observed topology of the 1.2 Jy survey to the predictions of dynamically evolved, unbiased, gravitational instability models that have Gaussian initial conditions. The model with an $n=-1$, power-law initial power spectrum achieves the best overall agreement with the data, though models with a low-density cold dark matter power spectrum and an $n=0$ power-law spectrum are also consistent. The observed topology is inconsistent with an initially Gaussian model that has $n=-2$, and it is strongly inconsistent with a Voronoi foam model, which has a non-Gaussian, bubble topology. ", + "introduction": "According to the most popular theories of structure formation, the observed distribution of galaxies --- a complex network of clusters, superclusters, tunnels, and voids --- developed by gravitational instability from Gaussian primordial fluctuations. Two different and complementary approaches have been followed to test the Gaussian hypothesis. The first uses the probability distribution function (PDF) or its moments (skewness, kurtosis, etc.); observed results are compared to predictions for a gravitationally evolved Gaussian field, which are computed either by numerical simulations or by various approximation schemes (e.g., Fry 1984; Bernardeau 1992; Juszkiewicz, Bouchet, \\& Colombi 1993; Juszkiewicz et al. 1995; Bernardeau \\& Kofman 1995; Protogeros \\& Scherrer 1997). The second approach uses topological characteristics of the galaxy density field, quantified by percolation analysis (Shandarin \\& Zel'dovich 1983; Yess, Shandarin \\& Fisher 1997) or by the genus of isodensity contours (Gott, Melott, \\& Dickinson 1986, hereafter GMD; Gott, Weinberg, \\& Melott 1987, hereafter GWM; for a review see Melott 1990). The genus measure also yields constraints on the index of the primordial power spectrum by quantifying the ``corrugation'' of structure in the smoothed density field. In this paper we apply the genus method to one of the largest complete galaxy redshift surveys, the 1.2 Jy IRAS redshift survey (Fisher et al. 1995). We make extensive use of mock catalogs drawn from cosmological N-body simulations to estimate systematic and random errors and to evaluate the viability of models. Our statistical methodology should also be useful for topological analyses of future, larger redshift surveys. Our basic approach follows that of GMD, GWM, and Gott et al. (1989). From the galaxy distribution, we create a density field by convolving with a Gaussian window, \\begin{equation} W(r)={1\\over {(2 \\pi)^{3/2}\\lambda^3}}e^{-r^2/{2\\lambda^2}}. \\label{window} \\end{equation} [Note that our definition of the smoothing length, $\\lambda$, based on the conventional form of a Gaussian window, differs by a factor of $2^{1/2}$ from that used in the Gott et al. papers.] We then construct isodensity contours at a variety of threshold levels and measure the genus of each. Applying the Gauss-Bonnet theorem, GMD define the genus $G_s$ through the integrated Gaussian curvature, \\begin{equation} G_s\\equiv -{1\\over {4\\pi}} \\int_S K dA, \\label{curvature} \\end{equation} with $K \\equiv 1/(a_1 a_2)$ where $a_1$ and $a_2$ are the principal radii of curvature. Equation~(\\ref{curvature}) differs slightly from the standard mathematical definition of the genus, but it is useful for cosmological purposes because it can be applied to a contour that runs into the boundary of a finite survey and it defines a quantity that is, statistically, proportional to volume. For a compact surface, $1+G_s$ is the number of handles or holes (in the sense of donut holes), while a surface broken into $n$ disjoint, simply connected pieces (e.g., $n$ spheres) has $G_s=-n$. We measure $G_s$ using the program CONTOUR (Weinberg 1988), based on the algorithm of GMD (see Coles, Davies, \\& Pearson 1996 for an alternative method of computing the genus). For a Gaussian random field, the mean genus per unit volume is \\begin{equation} g_s=A(1-\\nu^2)e^{-\\nu^2/ 2}, \\label{wcurve} \\end{equation} where $\\nu$, is the threshold density of the contour in units of the standard deviation (Doroshkevich 1970; Adler 1981; Bardeen et al. 1986; Hamilton, Gott, \\& Weinberg 1986). Positive and negative fluctuations are statistically interchangeable in a Gaussian field, so the $\\nu=0$ contour has a spongelike topology (positive $G_s$) in which the high and low density regions are both multiply connected and mutually interlocking. At high or low $\\nu$ the genus becomes negative as a typical contour breaks into separate bags around isolated clusters or voids, but the dependence is symmetric about $\\nu=0$. The normalizing constant $A$ depends on the second moment of the power spectrum. For a field with a power-law spectrum $P(k) \\propto k^n$ smoothed with the Gaussian filter of equation~(\\ref{window}), it is \\begin{equation} A= {1 \\over 4\\pi^2\\lambda^3} \\left( { {3+n}\\over{6}} \\right)^{3/2}. \\label{amplitude} \\end{equation} A field with more small scale power (higher $n$) has choppier, more corrugated structure, and hence a higher genus per unit volume for a given smoothing length. Since the smoothing length provides the only characteristic length scale in a Gaussian field with a power-law spectrum, the mean genus per unit volume necessarily scales as $\\lambda^{-3}$ for fixed $n$. Linear evolution preserves the Gaussian character of the initial density field. Nonlinear evolution does not, but the effects of nonlinear evolution on the genus curve are modest if the smoothing length is greater than or equal to the correlation length {\\it and} one characterizes contours by the fractional volume they enclose rather than the density level {\\it per se} (GWM; Weinberg, Gott \\& Melott 1987). Volume weighting also makes the genus curve insensitive to ``biased'' galaxy formation, since even nonlinear bias tends to maintain a monotonic relation between galaxy density and mass density. The information ``lost'' by volume weighting is precisely that contained in the PDF, so with this approach the genus curve and PDF at a given smoothing scale provide independent and complementary information about the density field. For convenience, we characterize a contour that encloses fractional volume $f$ (in the region above the threshold density) by the value of $\\nu$ for a corresponding contour in a Gaussian field, defined through the implicit equation \\begin{equation} f = (2\\pi)^{-1/2} \\int_\\nu^\\infty e^{-x^2/2} dx. \\label{nudef} \\end{equation} With this definition, equation~(\\ref{wcurve}) continues to hold for a Gaussian field, and it remains a good first approximation as the field evolves into the nonlinear regime (Melott, Weinberg, \\& Gott 1988; Park \\& Gott 1991). Using second-order perturbation theory, Matsubara (1994) and Matsubara \\& Suto (1996) show that even weakly non-linear evolution distorts the shape of the genus curve if it is plotted as a function of density contrast rather than fractional volume or equivalent $\\nu$. The genus statistic has been applied previously to six different redshift surveys of optically selected galaxies (Gott et al. 1989; Park, Gott, \\& da Costa 1992; Vogeley et al. 1994), and to the QDOT survey, a 1-in-6 subset of IRAS galaxies with $60\\mu$ flux density brighter than 0.6 Jy (Moore et al. 1992). It has also been applied to redshift surveys of Abell clusters (Gott et al. 1989; Rhoads, Gott, \\& Postman 1994), and its 2-dimensional analog has been applied to redshift slices (Park et al. 1992; Colley 1996), to projected galaxy and cluster catalogs (Coles \\& Plionis 1991; Plionis, Valdardini, \\& Coles 1992; Gott et al. 1992), and to COBE maps of cosmic microwave background fluctuations (Smoot et al. 1994; Colley, Gott, \\& Park 1996; Kogut et al. 1996). The 1.2 Jy IRAS survey contains 5321 galaxies and covers all of the sky except for the Galactic plane (Galactic latitude $|b|<5^\\circ$) and a number of small, isolated patches at high Galactic latitude, covering about 4\\% of the sky in total. For our analysis we use a catalog provided by M.\\ Strauss in which these high latitude regions have been filled with randomly placed galaxies. Observational details of the survey are described in Strauss et al. (1992) and Fisher et al. (1995). The survey has been the basis for many statistical investigations of large scale structure including the power spectrum (Fisher et al. 1993; Cole, Fisher, \\& Weinberg 1995), the two-point correlation function (Fisher et al. 1994ab), and moments of the counts-in-cells distribution (Bouchet et al. 1993). It has also been used in comparisons between predicted and observed peculiar velocity fields (e.g., Davis, Nusser, \\& Willick 1996). For topological analysis we use volume limited subsets of the redshift survey, so that the physical properties of the tracer galaxies and the effects of shot noise are uniform throughout the survey volume. A volume limited sample consists of those galaxies within distance $\\rmax$ that are luminous enough that they would still exceed the survey flux limit if they were at distance $\\rmax$. The number of galaxies in a volume limited sample first increases with $\\rmax$ as the survey volume grows, then declines at large $\\rmax$ as the fraction of galaxies luminous enough to be seen at the sample edge begins to decline rapidly. For the 1.2 Jy survey, the size of a volume limited sample peaks at $\\sim 1100$ galaxies for $\\rmax \\approx 60\\hmpc$.$^1$ \\footnotetext[1] {$h \\equiv H_0 / 100 \\,\\mkms\\,{\\rm Mpc}^{-1}$.} We infer a galaxy's distance from its redshift referred to the frame of the Local Group. The smoothing length must be large enough to suppress shot noise fluctuations in the density field, but at fixed $\\rmax$ increasing the smoothing length reduces the number of independent resolution elements in the survey volume. Following the rule of thumb suggested by Weinberg et al. (1987) and used in subsequent observational analyses, we adopt a smoothing length $\\lambda \\approx \\dbar/\\sqrt{2}$, where $\\dbar \\equiv n_g^{-1/3}$ is the mean intergalaxy separation. (The $\\sqrt{2}$ factor does not appear in the earlier papers because of their different Gaussian filter definition.) We discuss this choice further in \\S 3 below. The number of resolution elements in the smoothed density field, i.e., the ratio of the survey volume to the smoothing volume, is \\begin{equation} \\nres= {\\omega_s \\rmax^3 \\over 3}\\, {1 \\over (2 \\pi)^{3/2} \\lambda^3}, \\label{nres} \\end{equation} where $\\omega_s = 4\\pi(1-\\sin{\\pi \\over 36})$ sr is the solid angle of the 1.2 Jy survey. \\begin{figure} \\centerline{ \\epsfxsize=3.5truein \\epsfbox[105 445 515 680]{f1.eps} } \\caption{\\protect\\small The mean separation $\\bar d$ (solid line, left axis scale) and the number of resolution elements $\\nres$ (dotted line, right axis scale) for volume limited samples of depth $\\rmax$. Distances are in $\\hmpc$. \\vskip 2.0truein \\null \\label{fig:nres} \\normalsize} \\end{figure} Figure~1 plots the mean separation $\\dbar$ and the number of resolution elements $\\nres$, computed from equation~[\\ref{nres}] assuming $\\lambda=\\dbar/\\sqrt{2}$, as a function of the sample depth $\\rmax$. The number of resolution elements (which is proportional to the number of galaxies in the volume limited sample) peaks at $\\sim 170$ for $\\rmax=60\\hmpc$, where the mean separation is $\\dbar=9.5\\hmpc$. In this sense, the $\\rmax=60\\hmpc$ sample is the optimal volume limited subset that we can construct for topological analysis, and we focus our greatest attention on this sample. The corresponding smoothing scale is $\\lambda \\approx 7\\hmpc$. Since the dependence of topology on smoothing scale is itself interesting, we also analyze samples with $\\rmax=30\\hmpc$ and $100\\hmpc$, using smoothing scales $\\lambda = 4\\hmpc$ and $12\\hmpc$, respectively. In the next section we describe our procedure for creating mock catalogs designed to mimic the 1.2 Jy redshift survey. In \\S 3 we use these mock catalogs to study the systematic distortions in the genus curve that arise from shot noise, the finite size of the survey volume, and peculiar velocities. In \\S 4 we use the mock catalogs to examine the magnitude and covariance of random errors in the genus curves of the 1.2 Jy subsamples, and we outline a statistical methodology for comparing observed and predicted genus curves. In \\S 5 we present results for the 1.2 Jy survey and compare them to theoretical predictions. We summarize our conclusions in \\S 6. ", + "conclusions": "We have measured the topology of the galaxy distribution in the IRAS 1.2 Jy redshift survey using the methods of GMD, GWM, and Gott et al. (1989). We consider three volume limited subsets of the data, with limiting radii $\\rmax=30$, 60, and $100\\hmpc$, analyzed with corresponding smoothing lengths $\\lambda=4$, 7, and $12\\hmpc$. We use mock catalogs drawn from cosmological N-body simulations in order to derive theoretically predicted genus curves and to study the systematic and random errors expected in samples of this size. Our principal conclusions are: \\\\ (1) In tests on mock catalogs from low-$\\Omega$ CDM simulations, the net systematic error in volumes the size of our 1.2 Jy subsamples is small compared to the random errors. However, this small net error involves a cancellation between the systematic effects of measuring the genus in a volume that contains few independent structures and the effects of smoothing only with that portion of the smoothing window that lies within the sample volume. \\\\ (2) The covariance matrix of random errors in the genus curve is predominantly diagonal, but there are significant correlations in the errors that should be taken into account when assessing theoretical models. To a reasonable approximation, $\\chi^2$ (including covariances) is distributed as it would be if the errors followed a multivariate Gaussian distribution, though extreme values of $\\chi^2$ are more common than they would be for purely Gaussian statistics. \\\\ (3) With $\\lambda=12\\hmpc$, the genus curve of the 1.2 Jy data has a symmetric form similar to that predicted for a Gaussian random field. For $\\lambda=7\\hmpc$ and $4\\hmpc$, the observed genus curves are increasingly shifted in the direction of a ``meatball'' topology. \\\\ (4) Taken individually, the three observed genus curves are consistent at the $>5\\%$ level with the topology of dynamically evolved N-body models that have Gaussian initial conditions with low-$\\Omega$ CDM ($\\Gamma=0.25$), $n=0$, or $n=-1$ power spectra. Combining all three data sets, the $n=-1$ model is the most successful overall, with a likelihood ratio of 91.8 relative to CDM and 161 relative to $n=0$. \\\\ (5) The observed genus curves are inconsistent with an $n=-2$, initially Gaussian model, which produces structure that is excessively coherent and, consequently, genus curves whose amplitudes are too low. The observed genus curves are strongly inconsistent with a Voronoi foam model, which, on account of its ``bubble'' topology, predicts genus curves that are systematically shifted towards higher densities, in the opposite direction from the observed shifts. Our conclusions about the shapes of observed genus curves and their consistency or inconsistency with various theoretical models are similar to those drawn from a number of other topological studies of the galaxy distribution (Gott et al.\\ 1989; Moore et al.\\ 1992; Park et al.\\ 1992; Colley 1996). They are somewhat at odds with the results of Vogeley et al. (1994), who found that genus curves measured from the extended CfA redshift survey showed shifts in the direction of a bubble topology and were inconsistent with the genus curves predicted by CDM N-body models. The differences could reflect systematic differences in the structure traced by optical and IRAS galaxies, differences in the details of the topology analysis, or the somewhat greater statistical power of the CfA data set at the short smoothing lengths ($\\lambda \\sim 5\\hmpc$) where the differences are most pronounced. Earlier topology studies, recognizing the problem of correlated errors in the genus curve, have developed ``meta-statistics'' that characterize the overall shape of the genus curve (e.g., amplitude, asymmetry, width) and used these to assess the compatibility of models with the observations (see, e.g., Vogeley et al.\\ 1994). Here we have taken the more direct approach of measuring the error covariance matrix from mock catalogs and incorporating it into model assessments. We compute $\\chi^2$ values that include the error covariance (eq.~[\\ref{chisqr}]) and calculate the distribution of $\\chi^2$ from mock catalogs, in order to get an absolute, frequentist assessment of a model's goodness-of-fit. The distribution of $\\chi^2$ values in the mock catalogs implies that the multivariate Gaussian approximation describes the error distribution quite well, failing only in the extreme tails. We perform likelihood ratio comparisons between theoretical models to assess their relative ability to account for the observed topology data, making use of the Gaussian-error approximation. The advantages of our approach are that it has a clear statistical motivation, it provides a natural path for combining information from independent data samples, and, to the extent that the Gaussian-error approximation holds, it makes the best possible use of the data because it is based directly on the likelihood. The disadvantage is that a high $\\chi^2$ or low likelihood value says nothing in itself about how the model prediction and the data disagree. Thus, the likelihood approach used here is the most statistically powerful way to assess and compare models, but measures like those of Park et al.\\ (1992) and Vogeley et al.\\ (1994) may be useful for quantifying the nature of discrepancies between theory and observation. Our approach to handling correlated errors is similar to that used by Fisher et al.\\ (1994) and Cole et al.\\ (1995) in studies of redshift space distortions of the correlation function and power spectrum, and it can be adapted to many other problems in which error correlations are important but computable. It should be especially useful for topological analyses of future large galaxy redshift surveys, such as the Anglo-Australian 2dF survey and the Sloan Digital Sky Survey, which will yield much more stringent tests of the hypothesis that structure in the universe formed from Gaussian primordial fluctuations." + }, + "9701/astro-ph9701057_arXiv.txt": { + "abstract": "We report on a deep search with the Westerbork Synthesis Radio Telescope towards the galactic anticenter for the 327~MHz hyperfine transition of DI. This is a favorable direction for a search because:~(i)~the HI optical depth is high due to velocity crowding; (ii)~the observed molecular column density is low (implying that most of the deuterium would probably be in atomic form, rather than in HD); and (iii)~the stellar reprocessing should be minimal. Our observations are about a factor of two more sensitive than previous searches for DI in this direction. We detect a low significance ($\\sim4~\\sigma$) feature, consistent in both amplitude and center frequency with an emission feature reported previously (Blitz \\& Heiles 1987). If this is the DI line, then the implied $N_{\\rm D}/N_{\\rm H}$ of $3.9\\pm1.0\\times10^{-5}$ is comparable to the inferred pre-solar deuterium abundance. Our observation is consistent with the recent low measurements of D/H towards high-redshift Lyman-limit systems. On the other hand, if the reports of high DI abundance ($\\sim 24\\times10^{-5}$) in such systems are confirmed, then our observations imply that even in regions of reduced star formation within the outer Galaxy, the DI abundance has been reduced by a factor of $\\sim6$ from the primordial abundance. ", + "introduction": "In the standard big bang model, the primordial abundance of deuterium is a sensitive function of the baryon to photon ratio, (e.g. Walker et al. 1991) making it a quantity of great cosmological interest. Further, since all known astrophysical processes (apart from the big bang itself, of course) result in a net destruction of deuterium, the currently observed value of the deuterium abundance is a strict lower limit to its primordial abundance. HST observations of the DI Lyman$-\\alpha$ line in the local solar neighborhood (Linsky et al. 1993) yield a deuterium abundance of $N_{\\rm D}/N_{\\rm H} = 1.65^{+0.07}_{-0.18}\\times 10^{-5}$. Conversion from this local current abundance to the primordial abundance depends on less well understood details of the history of the stellar reprocessing of matter in the local ISM. In order to circumvent this problem a number of groups have been attempting to measure the deuterium abundance in high-redshift Lyman-limit systems. Since these systems are less chemically evolved than the local ISM, conversion from the measured to the primordial abundance should be more straightforward. However, the results of these observations are conflicting, with different groups (Carswell et al. 1994, Songalia et al. 1994, Tytler et al. 1996) measuring abundances which differ by more than an order of magnitude. Part of the problem is that the DI Lyman$-\\alpha$ line is separated from the HI Lyman$-\\alpha$ line by only 82~km~s$^{-1}$, making the chance of contamination of the DI line by absorption from a small parcel of HI at a slightly different velocity from that of the main Lyman-limit system non-negligible. There have also been several attempts to observe the hyperfine transition of DI at radio frequencies. Since the frequency of this line is more than a factor of 4 lower than the frequency of the corresponding transition in HI, the question of HI contamination does not arise. Two kinds of lines-of-sight have been favored in the past, the first towards bright radio sources (Sgr~A \\& Cas~A; Weinreb 1962, Anantharamiah \\& Radhakrishnan 1979, Heiles et al. 1993), where the hydrogen column density is known to be high. The disadvantages of these lines of sight are that: (i)~the bright radio sources contribute significantly to the system temperature, making detection more difficult; (ii)~any measurement refers only to the thin pencil beam subtended by the absorber; and (iii)~the molecular column density is also high, making it likely that most of the deuterium is in molecular rather than atomic form (Heiles et al. 1993). Anantharamiah \\& Radhakrishnan (1979) placed an upper limit of $5.8\\times 10^{-5}$ on the DI abundance towards Sgr~A. Heiles et al. (1993) reached similar limits towards Sgr~A as well as Cas~A. The other promising direction for a search for the radio emission from DI is that towards the galactic anticenter, where one expects the line to be in emission. The advantages of this direction are that (i)~the high optical depth of HI is due to velocity crowding along a long pathlength rather than a high volume density; (ii)~the molecular column density and metallicity are low; and (iii)~the observations are sensitive to the DI abundance within the entire telescope beam, and not just a narrow cone towards the background source as in the case of absorption observations. The results of a long integration in the direction $(l,b)~=~(183\\fdg 0,+0\\fdg 5)$ using the Hat Creek telescope were presented by Blitz \\& Heiles (1987), who found an upper limit ($2~\\sigma$) of $6.0\\times 10^{-5}$ for $N_{\\rm D}/N_{\\rm H}$. Here we report on a long integration towards a partially overlapping line-of-sight, $(l,b)~=~(183\\fdg 0,-0\\fdg 5)$, with the Westerbork synthesis array. Using the 14 WSRT telescopes as independent single dishes allowed us to significantly increase the effective integration time. ", + "conclusions": "The galactic anticenter region, near $(l,b)~=~(183\\fdg 0,0\\degr )$, is particularly well-suited to a search for the DI hyperfine emission feature at $\\lambda$92-cm. The HI emission brightness in this region peaks smoothly into a plateau extending spatially over several degrees with a brightness temperature 135$\\pm$3~K, as can be seen in Burton (1985). Conditions should therefore be reasonably constant over the $2\\fdg 5$ beam of the WSRT dishes. The extended region of line brightening is paired with a line narrowing due to velocity crowding. The HI emission-line profile from this region is relatively narrow, with a FWHM of 18.7~km~s$^{-1}$. The line core is rather flat-topped, strongly suggesting a high line opacity. It also displays shallow self-absorption, indicating a modest amount of temperature substructure along the line-of-sight. Based on the depth of the self-absorption features and assuming a substantial line opacity, the spin temperature of the atomic gas appears to be in the range 125--135~K. Direct measurements of the HI opacity along many lines of sight in this region have not yet been carried out. The typical sky density of moderately bright extragalactic background sources suitable for such a measurement suggests that a few tens of lines of sight could, in principle, be observed. The closest line of sight that has been observed in HI absorption (Dickey et al. 1983) is at $(l,b)~=~(189\\fdg 6,-0\\fdg 6)$. This is near the edge of the plateau in emission brightness, where the line profile is already somewhat broader than in the direction of maximum velocity crowding, with a FWHM of 23.4~km~s$^{-1}$. The absorption profile extends smoothly over the entire velocity extent in which the emission profile exceeds a brightness of about 5--10~K, with a uniform high opacity of about 2 across the line core. This is consistent with the expectation for a gas distribution in which the column density is dominated by an approximately isothermal gas with kinetic temperature in the range 125--135~K. An absorption equivalent width, $\\int\\tau~{\\rm d}V$~=~33.68~km~s$^{-1}$ is found. The very smooth nature of the HI emission suggests that it is plausible to expect a comparable equivalent width in the direction $(l,b)~=~(183\\fdg 0,-0\\fdg 5)$. Assuming that the DI emission is either mixed with, or lies behind, most of the galactic continuum emission, (which has the substantial brightness of about 70~K at 92~cm wavelength) the equation of radiative transfer yields simply: \\begin{equation} T_{\\rm B} = T_{\\rm S}(1- e^{-\\tau_{\\rm D}}) \\end{equation} where $T_{\\rm B}$ is the differential brightness temperature on and off the line, $T_{\\rm S}$ is the spin temperature of the DI and ${\\tau_{\\rm D}}$ is the optical depth of the DI. For ${\\tau_{\\rm D}}~<<~1$ this yields $T_{\\rm B}~=~T_{\\rm S}\\tau_{\\rm D}$, or $\\int~T_{\\rm B}~{\\rm d}V~=~T_{\\rm S}\\int\\tau_{\\rm D} {\\rm d}V$ for an approximately isothermal gas. The line integral of the Gaussian overlaid on the possible feature in Figure~1 is $\\int~T_{\\rm B}~{\\rm d}V~=~0.048\\pm0.012$~K~km~s$^{-1}$. Assuming that the spin temperature of the DI is the same as that of HI (i.e. 130~K), this corresponds to $\\int\\tau_{\\rm D}~{\\rm d}V$~=~3.7$\\pm0.9\\times$10$^{-4}$~km~s$^{-1}$. Hence the estimate of the ratio of the optical depths is ${\\tau_{\\rm D}/\\tau_{\\rm H}}~=~1.1\\pm0.3 \\times 10^{-5}$. The relationship between the ratio of the optical depths and the ratio of the column densities of DI and HI is (Anantharamiah \\& Radhakrishnan 1979) ${\\tau_{\\rm D}/\\tau_{\\rm H}}~=~0.28~{N_{\\rm D}/N_{\\rm H}}$, hence the estimated abundance of DI is $3.9\\pm1.0\\times10^{-5}$. This is comparable to the inferred pre-solar abundance (Kunde et al. 1982, Courtin et al. 1984), and about a factor of 2 above the current DI abundance in the local solar neighborhood (Linsky et al. 1993). As discussed in the introduction, the lack of detection of DI towards the inner Galaxy suggests that most of the DI has been converted into HD within molecular clouds (Heiles et al. 1993). Towards the outer Galaxy, however, the molecular column density is low along many sight lines. In fact, the survey by Dame et al. (1987) detected very little CO emission in the galactic plane between 180 and 185 degrees longitude except for a small concentration at $-10$~km~s$^{-1}$. The high optical depth in HI seems to be due to velocity crowding in diffuse gas along a pathlength of several kpc, and does not appear to arise in a single physical entity like a giant molecular cloud. It seems plausible, therefore, that a large fraction of the deuterium in this direction is still in atomic form. Further, the metallicity gradient in the Galaxy suggests that the DI abundance in the outer Galaxy would be higher than the inner. Calculations by Prantzos (1996) show that the deuterium abundance gradient is in general steeper than (and of course opposite in sign to) the oxygen gradient (because of the late ejection of metal-poor but deuterium-free gas from low mass stars formed at early times). While the exact gradient is sensitive to the assumed infall model, the current DI abundance within diffuse atomic gas in the inner and outer Galaxy could well differ by a factor of two. Another useful comparison to make is with the range of measured DI abundances, $2.3-24\\times10^{-5}$, in high redshift Lyman limit systems (Carswell et al. 1994, Songalia et al. 1994, Tytler et al. 1996). If a high primordial abundance turns out to be correct, our results indicate that even in regions with reduced star formation, the current deuterium abundance is about a factor of 6 lower than primordial. In contrast, recent model calculations of astration of deuterium suggest that the abundance evolution is modest, with current abundance only a factor $\\sim2$ less than primordial (Galli et al. 1995). On the other hand, the low value of D/H measured by Tytler et al. (1996), $2.3\\pm0.3\\times10^{-5}$, is consistent with our own possible detection and the inferred pre-solar value. Direct imaging of DI emission in the outer regions of nearby galaxies may provide a very effective means of addressing this issue comprehensively, once the capability for achieving sub-mK sensitivity on arcmin scales becomes available with the next generation of radio telescopes. The Giant Meterwave Radio Telescope (GMRT) should already allow a robust detection of DI emission from the Galaxy within the next few years, while DI imaging of nearby external galaxies should become possible with construction of the Square Kilometer Array Interferometer (SKAI) early in the mext millenium." + }, + "9701/astro-ph9701105_arXiv.txt": { + "abstract": "{\\vspace*{27pt}ABSTRACT\\par\\relax} \\def The International AGN Watch has monitored a number of radio-quiet and radio-loud Active Galactic Nuclei -- the most luminous objects in the universe. We present a review of the main observational results from the continuum monitoring campaigns, concentrating on those which aim to quantify the simultaneous ultraviolet to X-ray variability characteristics. These data provide strong constraints on the proposed continuum emission mechanisms. The AGN Watch campaigns have made extensive use of a wide variety of both ground- and space-based multi-waveband observational facilities, and we stress that long-term simultaneous access to the entire electromagnetic spectrum is essential if further progress is to be made. ", + "introduction": " ", + "conclusions": "" + }, + "9701/astro-ph9701043_arXiv.txt": { + "abstract": "We study metal line absorption of CIV, CII, SiIV, and NV at redshifts $z=3.5$ to $z=2$ within the framework of a cosmological model for the Lyman alpha forest, comparing the results of numerical simulations to recent observations by Songaila \\& Cowie (1996, SC). In agreement with Rauch, Haehnelt \\& Steinmetz (1996), we find that the observed mean value of the CIV/HI ratio at $z \\simeq 3$ is reproduced if a uniform metallicity of $[{\\rm C}/{\\rm H}]\\sim -2.5$ is assumed in our model, but that the observed scatter in this ratio is larger than predicted, implying a scatter in the metallicity of the absorbing systems of roughly an order of magnitude. We further argue that absorbers with relatively low column densities ($\\log{N_{\\rm HI}} < 15$) likely have a mean metallicity [C/H] less than $-2.5$, a result which is basically independent of the model considered. The enrichment pattern that is required for our model to match SC's CIV observations is very similar to that {\\it predicted} by Gnedin \\& Ostriker's (1997) simulations of reionization and metal enrichment by population III stars in this type of cosmological scenario. Our model predicts no significant evolution in the mean values of metal line column densities between $z=3.5$ and $z=2$. Comparison of the predicted and observed numbers of SiIV and NV systems suggests that the photoionizing background radiation field at $z\\sim 3$ is somewhat softer than that proposed by Haardt \\& Madau (1996). Our model does not account for the increase in the SiIV/CIV ratio at $z \\simgt 3.2$ found by SC. While SC suggested that the increase could be explained by a softening of the radiation field due to HeII absorption at $z \\simgt 3$, such a modification does not significantly raise the mean value of SiIV/CIV in our simulation because it shifts numerous weak SiIV systems to just above the detection limit, thus keeping the mean column density of observable SiIV systems low. ", + "introduction": "Spectra of high redshift quasi-stellar objects (QSOs) are densely populated with absorption lines blueward of the Lyman alpha emission peak. This so-called ``Lyman alpha forest'' is interpreted as arising from absorption by neutral hydrogen cosmologically distributed along the line-of-sight to the QSO (e.g. Lynds 1971; Sargent \\etal 1980). For more than a decade it was not known whether the gaseous systems responsible for the Lyman alpha forest were chemically pristine. Meyer \\& York (1987) presented early evidence for metal lines in the forest, and within the last several years observations with improved spectral resolution and signal-to-noise ratio (S/N) have revealed the presence of heavy elements in a significant fraction of Lyman alpha systems with moderate HI column densities ($\\sim 10^{14}-10^{17} {\\rm cm}^{-2}$). Analyzing data from the HIRES spectrograph on the Keck 10~m telescope, Cowie \\etal (1995) find that $\\sim 50$\\% of identified Lyman alpha lines with $\\log{{\\rm N}_{\\rm HI}} > 14.5 $ have associated CIV absorption. This finding is roughly consistent with that of Womble \\etal (1995) who infer CIV absorption in 40--45\\% of all Lyman alpha systems with $\\log{{\\rm N}_{\\rm HI}} > 14.3$. More recently, Songaila \\& Cowie (1996, hereafter SC) presented an analysis of even higher S/N spectra and found CIV absorption in $\\simeq 75\\%$ of all systems with $\\log{{\\rm N}_{\\rm HI}} > 14.5$ and in $\\simeq 90-100$ \\% of those with $\\log{{\\rm N}_{\\rm HI}} > 15.2$. In addition, they detected absorption from SiIV, CII, and NV in a number of these systems. In the present paper, we analyze the predicted metal line absorption from a hydrodynamic cosmological simulation of the high-redshift intergalactic medium, and compare these predictions to the observations of SC. The analysis of such simulations during recent years (e.g. Cen \\etal 1994; Zhang \\etal 1995; Petitjean \\etal 1995; Hernquist \\etal 1996) has established a convincing general scenario in which the Lyman alpha forest is produced by regions of moderate overdensity in a hierarchically clustering universe, pervaded by an ionizing radiation field presumed to originate from the QSOs (for related semi-analytic modeling see, e.g., Bi 1993; Bi \\& Davidsen 1996; Hui, Gnedin, \\& Zhang 1997). These cosmological Lyman alpha forest models predict QSO absorption features that resemble those observed, and they reproduce the column density and Doppler parameter distributions down to the lowest observable column densities (Miralda-Escud\\'e \\etal 1996; Dav\\'e \\etal 1997, hereafter DHWK; Zhang \\etal 1997). The interpretation of the metal line data within the framework of these models of the Lyman alpha forest constitutes a further test of the validity of the models and provides potentially valuable insight to the physical conditions prevailing in the intergalactic medium, such as gas densities, temperatures, and the shape and intensity of the ionizing radiation field. It is even conceivable that comparisons between such models and the wealth of forthcoming observations will constrain the underlying cosmological parameters, although this possibility remains to be demonstrated. Haehnelt, Steinmetz \\& Rauch (1996) and Rauch, Haehnelt \\& Steinmetz (1996) have pioneered the technique for incorporating metals into cosmological Lyman alpha models. They endow the gas with a uniform metallicity, postulated to have arisen from a population III burst of star formation at a significantly earlier time, and then compute the ionization states based on the assumption of a uniform ionizing background of a given strength and spectral shape. By generating and analyzing artificial metal line spectra along lines of sight through cosmological hydrodynamical simulations, these authors compare their models to a number of observables such as the Doppler parameter and column density distributions, CIV/HI and other metal line ratios, and two point correlation functions, and generally find acceptable agreement between their models and observations. In the present work we adopt a similar approach, but whereas the gaseous regions studied in the above mentioned works are overdense regions which eventually evolve into galaxies, we utilize a larger simulation which allows us to consider a random region of the universe and make quantitative predictions regarding the statistics of observable metal systems and their evolution in redshift. Through most of this paper, we will assume that the simulation provides an adequate physical model of the intergalactic medium and will use the comparison to observations to obtain information about the metal enrichment of the diffuse gas and the spectral shape of the ionizing radiation field. This assumption is supported by the previous successes of the model mentioned above, and we will argue in Section 4 that the success of the model in predicting the relative numbers of observable metal line systems provides further evidence favoring this physical picture of the high-redshift intergalactic medium. ", + "conclusions": "We have compared a cosmological simulation of the Lyman alpha forest and associated metal lines to recent observations of CIV, CII, SiIV, and NV absorption in high-redshift QSO spectra. Our simulated spectra have signal-to-noise properties similar to those of the observed spectra, and we analyze them by a Voigt-profile fitting procedure that is similar to that used in the observational analyses. We can examine our results from two different points of view. First, we can assume that our theoretical model provides an accurate description of physical conditions in the absorbing gas and use the comparison to the observations to draw inferences about metal enrichment of the IGM and the spectral shape of the ionizing radiation background. Second, we can ask whether the predicted properties of the IGM are consistent with the metal line observations, within the uncertainties of the observations and the theoretical modeling. Adopting the first point of view, we conclude: \\begin{enumerate} \\item{The mean metallicity of the IGM at $z\\approx 3$ corresponds to [C/H]$\\approx -2.5$. More specifically, for the Lyman beta selected sample of SC ($\\log{N_{\\rm HI}} \\simgt 15.2$), the inferred mean metallicity ranges from [C/H]$\\approx -2.8$ to [C/H]$\\approx -2.3$, depending on the assumed spectral shape of the ionizing background.} \\item{Significant spatial variations of [C/H], with an rms scatter of roughly one order of magnitude, are required to match the observed scatter in CIV/HI.} \\item{There is a trend between mean enrichment and HI column density; Lyman alpha absorbers with low HI column densities (and correspondingly low physical densities) have systematically lower metallicities.} \\item{The ionizing background spectrum at $z \\approx 3$ must be somewhat softer than the spectrum predicted by HM, in order to explain the observed number of SiIV systems. However, if the spectrum is too much softer than the HM spectrum, there will be conflict with the observed HeII optical depth (see Croft et al.\\ 1996; Davidsen, Kriss, \\& Zheng 1996; Hogan et al.\\ 1996).} \\end{enumerate} The required enrichment of the IGM (mean value, scatter, and trend with density) is very similar to that predicted by Gnedin \\& Ostriker's (1997) numerical simulations of metal enrichment and reionization by population III stars. The required scatter in metallicity is also plausible on observational grounds, since it is similar to that seen in halo stars in our own galaxy (e.g. Ryan \\& Norris 1991). However, some of the observed scatter in CIV/HI could also be contributed by inhomogeneities in the ionizing background instead of metallicity variations. Within the current uncertainties of the observational data and the theoretical modeling, we expect the predictions presented here to be generic to cosmological models that have similar amounts of small-scale power at $z \\sim 2-3.5$. Several results of our comparison support the adopted physical model of the high-redshift IGM: \\begin{enumerate} \\item{With plausible assumptions about enrichment and the ionizing background, the model reproduces the observed abundances of CIV, CII, SiIV, and NV in Lyman alpha forest systems fairly well. This agreement suggests that the predicted densities and temperatures of the absorbing gas are at least approximately correct.} \\item{Consistent with SC's observations, the model predicts no substantial evolution in the mean value of $N_{\\rm CIV}$ for Lyman-beta-selected HI lines over the redshift range $2.5 \\simlt z \\simlt 3.5$.} \\item{The model successfully explains the observed trend towards lower CIV/HI ratios in high column density systems ($N_{\\rm HI} \\sim 10^{16}-10^{17} {\\rm cm}^{-2}$). This trend arises because of the correlation between column density and physical density, which causes the stronger absorbers to have more of their carbon in lower ionization states.} \\item{Qualitatively, the model reproduces observations that show multiple CIV components associated with strongly saturated Lyman alpha lines (see Figure~1). We will explore the model predictions more quantitatively in future work.} \\end{enumerate} The observation that is most difficult to explain in this model is the jump in the mean SiIV/CIV ratio at $z \\approx 3$, which SC suggest is due to a rapid change in the shape of the ionizing background spectrum associated with HeII reionization. In our model, there are many weak SiIV systems below the SC detection threshold. Softening the background spectrum moves these systems above the threshold, keeping the mean SiIV column density low. To the extent that the analyses overlap, our conclusions agree well with those of Haehnelt \\etal (1996) and Rauch \\etal (1996), despite differences in the numerical techniques, in the mass resolution of the simulations, and in the type of simulations (a single large box representing a randomly selected volume vs.\\ multiple, higher resolution simulations focused on collapsing regions). There are many potential improvements for future modeling, including higher resolution of the gas dynamics to better describe the structures giving rise to the moderate column density systems where metals are detected, a non-equilibrium treatment of the ionization of all species, and explicit incorporation of spatial variations in the metallicity of the gas and the intensity and shape of the background radiation field. Future studies can also examine alternative reionization scenarios and other cosmological models, to see which theories of cosmic structure formation can account for metal line absorption in the high-redshift universe." + }, + "9701/astro-ph9701039_arXiv.txt": { + "abstract": "We consider non-local effects which arise when radiation emitted at one radius of an accretion disk either heats or cools gas at other radii through Compton scattering. We discuss three situations: 1. Radiation from the inner regions of an advection-dominated flow Compton cooling gas at intermediate radii and Compton heating gas at large radii. 2. Soft radiation from an outer thin accretion disk Compton cooling a hot one- or two-temperature flow on the inside. 3. Soft radiation from an inner thin accretion disk Compton cooling hot gas in a surrounding one-temperature flow. We describe how previous results are modified by these non-local interactions. We find that Compton heating or cooling of the gas by the radiation emitted in the inner regions of a hot flow is not important. Likewise, Compton cooling by the soft photons from an outer thin disk is negligible when the transition from a cold to a hot flow occurs at a radius greater than some minimum $R_{\\rm tr,min}$. However, if the hot flow terminates at $R < R_{\\rm tr,min}$, non-local cooling is so strong that the hot gas is cooled to a thin disk configuration in a runaway process. In the case of a thin disk surrounded by a hot one-temperature flow, we find that Compton cooling by soft radiation dominates over local cooling in the hot gas for $\\dot{M} \\gsim 10^{-3} \\alpha \\dot{M}_{Edd}$, and $R \\lsim 10^4 R_{Schw}$. As a result, the maximum accretion rate for which an advection-dominated one-temperature solution exists, decreases by a factor of $\\sim 10$, compared to the value computed under an assumption of local energy balance. ", + "introduction": "Hot optically thin accretion flow solutions are often used to model X-ray binaries and active galactic nuclei. In addition to the original two-temperature model of Shapiro, Lightman \\& Eardley (1976) and the corona model of Haardt \\& Maraschi (1991), two new classes of advection-dominated models have been proposed recently: a two-temperature solution (Narayan \\& Yi 1995, Abramowicz et al. 1995), and a one-temperature solution (Esin et al. 1996). Advection-dominated models differ from the usual accretion solutions in that only a fraction of the viscously dissipated energy is radiated locally in the disk. The remainder of the generated energy is stored in the gas as entropy, resulting in a very hot, optically thin, quasispherical flow. If the accreting object is a black hole, as it is believed to be the case for several X-ray binaries and all active galactic nuclei, the stored energy is carried by the gas inside the horizon, so that the luminosity of the system can be orders of magnitude smaller than is predicted by standard accretion theory. The gas in advection-dominated solutions cools primarily through bremsstrahlung, synchrotron and inverse Compton scattering processes. Since the accreting electrons are heated to temperatures of up to $10^{10}\\,{\\rm K}$, the spectra of such systems span $\\sim 10$ orders of magnitude in energy, from synchrotron radio emission to $\\sim 400 keV$ X-rays produced via Compton cooling of hot electrons. Because two-temperature advection-dominated solutions have low radiative efficiency and hard power-law spectra, they have been successful in reproducing the observed properties of various low-luminosity X-ray sources including Sagittarius A$^*$ (Narayan, Yi, \\& Mahadevan 1995), the black hole nova A0620-00 in quiescence (Narayan, McClintock \\& Yi 1996), and the central source in NGC 4258 (Lasota et al. 1996). They also appear to be relevant to luminous systems (Narayan 1996). It is not yet clear if one-temperature solutions are relevant to any observed system. Though detailed numerical spectra of advection-dominated flows include global radiative transfer effects (Narayan, McClintock \\& Yi 1996; Narayan 1996; Narayan, Barret \\& McClintock 1997), analytical calculations of the physical properties of these accretion solutions, which we briefly review in \\S2, generally ignore the effects of non-local radiative transfer (Narayan \\& Yi 1995; Esin et al. 1996). This is a dangerous approximation. The optical depth for electron scattering in such flows is generally much lower than unity, and therefore, radiation emitted at one radius can potentially change the energy balance of the gas at quite different radii. It is the purpose of this paper to investigate the importance of this effect. We analyze three different but related issues. First, in the advection-dominated models the bulk of the luminosity is emitted from the innermost region of the disk ($R < 100 R_{Schw}$), where most of the gravitational energy is dissipated and the cooling efficiency of the gas is higher. The radiative flux from this region can have two effects on the rest of the flow, as we discuss in \\S3.1. The hot radiation can Compton heat (\\S3.1.1) the outer regions ($R > 1000 R_{Schw}$), where the gas is cooler (Shvartsman 1971, Ostriker at al. 1976, Grindlay 1978). On the other hand, cold but numerous synchrotron photons can contribute to cooling of the gas (\\S3.1.2) closer in ($100 R_{Schw} < R < 1000 R_{Schw}$). Next, the hot solutions based on the one-temperature advection model (Esin et al. 1996) do not extend down to the last stable orbit for $\\dot{M} > 0.003 \\alpha^2 \\dot{M}_{Edd}$. At accretion rates higher than this limit the only possible configuration is a standard thin disk (Shakura \\& Sunyaev 1973; Frank, King \\& Raine 1992), surrounded by a hot flow. Since most of the energy is dissipated in the thin disk, the number density of the cold photons emitted in the inner regions can be much greater than the number density of photons emitted locally by the hot gas, and we expect that the hot flow will suffer significant extra cooling via Compton scattering of the thin disk photons. This issue is discussed in \\S3.2. Finally, both one- and two-temperature advection-dominated solutions can be surrounded by a thin accretion disk, with the boundary between them at some radius $R_{tr}$, where the value of $R_{tr}$ depends on the parameters of the model. Here again the cold photons emitted by the thin disk can Compton cool electrons in the nearby hot flow. We analyze this situation in \\S3.3. Just as the radiation from the thin disk affects the energy balance of the hot flow, the cold gas itself is affected by irradiation from the advection-dominated zone. In \\S3.4 we investigate whether this process affects our conclusions from \\S3.2 and \\S3.3. We summarize and discuss the results in \\S4. ", + "conclusions": "In this paper we have explored how non-local radiation transfer affects the properties of hot one- and two-temperature advection-dominated solutions described by Esin et al. (1996), Narayan \\& Yi (1995) and Abramowicz et al. (1995). Since these solutions are optically thin, radiation emitted at one radius could in principle induce significant cooling or heating of the accreting gas at some other radius, through Compton scattering. We considered three different sources of non-local radiation: a luminous inner region of a hot disk ($R < 100 R_{Schw}$), a thin disk inside a hot one-temperature gas, and a thin Shakura \\& Sunyaev disk surrounding a hot advection-dominated flow. We found that two-temperature solutions are not affected by the radiation from their hot interior. These solutions are strongly advection-dominated everywhere except near the outer boundary of the disk, $R_{crit}$, (see Fig. 1[c]), and the inner part simply does not produce enough hot photons to induce significant heating of the outer layers. On the other hand, the gas density of the flow at $R > 100 R_{Schw}$ is so small, that with electron temperatures of $T_e \\lsim 10^9\\ {\\rm K}$, Compton cooling by non-local radiation can not play an important role either. The inner parts of one-temperature flows are cooling-dominated and produce more radiation than their two-temperature counterparts. However, single-temperature solutions are limited to lower values of the mass accretion rate (see Figure 1[a]), and since hot photons are produced mainly by Comptonization, which is not an efficient process at low $\\dot{M}$, no heating of the outer flows occurs. On the other hand, we find that Compton cooling by the photons from the interior dominates over local cooling processes for $10^2 R_{Schw}\\lsim R \\lsim 10^4R_{Schw}$ and $\\dot{M} \\gsim 10^{-7} \\dot{M}_{Edd}$. However, since the bulk of the observable radiation comes from inside $100 R_{Schw}$, the changes in the flow properties caused by extra cooling do not appreciably affect the overall spectrum of the system. Compton cooling by the radiation from the inner thin disk does have a significant impact on one-temperature flows, since it occurs at higher $\\dot{M}$, when Compton scattering is very effective. The extra cooling induced by the thin disk photons is so strong, that the maximum accretion rate for which the one-temperature model has an equilibrium solution decreases by a factor of $\\sim 10$ (Figure 3[a]). Finally, we found that the radiation from an outer thin disk does not have any effect on a hot flow that extends beyond $\\sim 10^3 R_{Schw}$. At this distance, the gas is too cold for significant Compton cooling to take place. When the transition radius between the hot and cold gas moves closer to the accreting object, the non-local cooling of two-temperature gas near the boundary can become very large. However, this effect is limited to the layers of the hot flow very close to the transition between the hot and cold gas; away from this boundary, the ratio of non-local to local cooling decreases as $r^{7/2}$. Therefore, external cooling does not affect the regions of the hot gas where most of the observable radiation is produced and can be disregarded until it is so strong that it forces the hot flow near the boundary to cool to a thin disk configuration. When this happens, the transition radius decreases and the hot gas in the new boundary layer experiences even stronger external cooling, which causes it to collapse in its turn. The resulting runaway process continues until the entire flow assumes the thin disk configuration. This phenomenon can be important if the hot advection-dominated flows are formed by evaporation of the thin disk near the accreting black hole (Narayan \\& Yi 1995). In that case, the maximum accretion rate at which such hot flow can form is in fact less than $\\dot{M}_{crit}$ computed based on the local assumption." + }, + "9701/astro-ph9701124_arXiv.txt": { + "abstract": "We present preliminary results on the effects of SNIa explosions on the Spectral Energy Distribution (SED) of distant galaxies and the possible modifications which may occur in integrated spectra, magnitudes and colours of simple galaxy models of different ages and metallicities few days after a SNIa event. ", + "introduction": "Recent observations allowed to derive spectra of very high redshift galaxies ($z \\gsim 3$) providing information on their earlier stage of evolution. Several authors infer galaxy ages by using the population synthesis technique and ages of the order of 1 Gyr or less have been suggested. We are interested in investigating the impact of SNe on the total emitted light of a galaxy in the range of ages running from 0.1 Gyr to few Gyr. The SNIa events should appear for ages comparable to the time scale of the first generation of white dwarf (WD) stars ($\\simeq 0.05 - 0.1$ Gyr). We suggest that SNIa may have a non--negligible impact on the SED of a galaxy. We assume few models for the stellar population synthesis which might be representative of the observed SED of high redshift galaxies and do not consider obscuration by dust. We stress that reproducing all the possible models which explain observed SEDs is not the goal of this work, while we focus our attention on the variations of the observable parameters when SNIa events are taken into account. ", + "conclusions": "" + }, + "9701/astro-ph9701062_arXiv.txt": { + "abstract": "We present CCD images of a sample of 39 HII galaxies taken at the Danish 1.54m telescope on La Silla. The images are used to analyse the morphology of these emission line dwarfs, and the structural properties of the knots of star formation and of the underlying galaxy. The sizes of the starbursts are measured. We propose a morphological classification based on the presence or absence of signs of tails, extensions, or distorted outer isophotes. This criterion segregates the objects into two broad morphological types with different physical properties: the more disturbed and extended (type I) HII galaxies having larger luminosities and velocity dispersions than the more compact and regular (type II) objects. The relative position of HII galaxies and of a sample of dwarf elliptical galaxies in the [R -- $\\sigma$] diagram support the hypothesis of a possible evolutionary link between the two types of galaxy. ", + "introduction": "HII galaxies are narrow emission line dwarf galaxies undergoing violent star formation (Melnick, Terlevich \\& Eggleton 1985) whose spectroscopic properties are indistinguishable from extragalactic giant HII regions in normal late type galaxies (e.g. 30 Dor in LMC, NGC 604 in M33) (Sargent \\& Searle 1970). Their high rates of star formation and low heavy element abundances imply that the star formation history must be simple and episodic (i.e. few burst of short duration followed by long quiescent periods). A recent review on the global properties of HII galaxies is given by Telles (1995, and references therein). The possible links of HII galaxies with other types of known dwarf galaxies have been discussed by Thuan (1983); Loose \\& Thuan (1985); Bothun \\etal (1986); Kunth, Maurogordato \\& Vigroux (1988); Davies \\& Phillipps (1988), Drinkwater \\& Hardy (1991). However, no conclusive answer has been given to the questions of what these systems will resemble when the present period of violent star formation ends, or what triggered the burst. It has been suggested that in their quiescent phase HII galaxies may be related to dwarf irregulars (dI) or dwarf elliptical galaxies (dE). Bothun \\etal (1986) made a comparative study of dIs and dEs in the Virgo cluster based on the colour distributions and structural properties derived from exponential fits to the surface brightness profiles (e.g. scale length and central brightness). They conclude that dIs are not progenitors of dEs, but they seem to form a parallel sequence of dwarf galaxies. The fading of dIs would make them very diffuse and place them below the detection threshold of photographic plates. They propose that Blue Compact Galaxies (BCG's, of which HII galaxies are a subset) could probably be gas-rich analog of dEs. Meurer, Mackie \\& Carignan (1994) have studied the structural properties of the dwarf amorphous galaxy NGC 2915 and compared with the properties of NGC 1705 (Meurer, Freeman \\& Dopita 1992) and NGC 5253 from the work of and S\\'ersic \\& Donzelli (1992). They find that their luminosity profiles show two components indicating the presence of two distinct stellar populations. The inner component represents the fraction of the galaxy dominated by hydrogen gas photoionized by the embedded massive star clusters. Its (B-R) colour profile is increasingly bluer inwards. The outer component has an exponential luminosity (also found for dE's and dIrr's) and a constant redder colour likely to represent an old stellar population remnant from a previous burst of star formation. Their main conclusion is that these galaxies are nearby BCG's that may provide a better insight on the properties of this type of galaxy and their connection with other dwarf galaxies. Kunth, Maurogordato \\& Vigroux (1988) analysed a small sample of BCG's to derive surface brightness profiles by ellipse fitting to different isophotal levels. Their results show that the BCG's present a \\,``mixed bag of morphologies\\,''. They find that the outer parts of the galaxies can be best fitted by a power law compared with those of elliptical galaxies. No definitive study has been made on the morphology of HII galaxies (or BCG's for that matter) up to now. The previous attempts have shown an extensive range of shapes from the most compact and apparently isolated to some clearly revealing diffuse extensions, multiple tails, and visually merging systems (Loose \\& Thuan 1985; Melnick 1987; Kunth, Maurogordato \\& Vigroux 1988; Salzer, MacAlpine \\& Boroson 1989b). Loose \\& Thuan (1985) have devised a classification scheme based on the shape and location of the burst in relation with the whole optical structure and the shapes of the outer envelopes. Melnick (1987) describes the systems in terms of being interacting, multiple, or isolated. He has preliminarily reported that 50\\% of the HII galaxies in his sample are star-like and isolated. Salzer, MacAlpine \\& Boroson (1989b), on the other hand, adopted a more detailed classification scheme. Primarily based on the absolute magnitude, size, and morphology, with some spectroscopic information as a secondary indicator, they identified 10 different classes of objects for a sample of emission line objects from the University of Michigan (UM) objective prism survey. Their sample includes some Seyfert galaxies as well as interacting pairs of disk galaxies, starburst nuclei and giant irregulars. Most HII galaxies in our sample are classified as \\,``dwarf HII hot spot galaxies\\,'', \\,``HII hot spot galaxies\\,'' or \\,``Sargent \\& Searle objects\\,''. We have used surface photometry in order to study the morphology and structural properties of HII galaxies. In Section~\\ref{morph:sample}, we first present the data sample. In \\S\\,\\ref{morph:results}, we present our results based on the analysis of the CCD images, as well as structural aspects based on the luminosity profiles. In \\S\\,\\ref{morph:discussion}, we discuss the results, and finally in \\S\\,\\ref{morph:conclusions} we present some conclusions. ", + "conclusions": "\\label{morph:conclusions} The morphological and structural properties of HII galaxies are studied in this paper. The main conclusions of this study are: \\begin{itemize} \\item HII galaxies can be classified within two broad morphological types: \\begin{description} \\item[Type I] irregular systems with signs of distorted outer isophotes, tails, wisps, fans, etc. \\item[Type II] regular and compact systems with symmetric morphology. \\end{description} Type I's were found to have higher luminosities and velocity dispersions than Type II's, while the equivalent widths of H$\\beta$ and oxygen abundances of the two types are roughly similar. This seems to indicate that the starbursts may have been triggered by different mechanisms in the two classes of objects. \\item We find three main types of light profiles in HII galaxies. The profile types {\\em qualitatively} relate to the overall morphology of the galaxies. The outer parts of the luminosity profiles of HII galaxies are well represented by an exponential scaling law. This will allow a direct comparison of the structural parameters (scale length and central surface brightness) with other types of dwarf galaxies and will lead us to derive important structural properties of the underlying systems once calibrated images are obtained. \\item While the burst sizes of HII galaxies are of the order of hundreds of parsec, the \\,``true\\,'' core radii of HII galaxies are basically unresolved and probably only few parsecs across. Yet, most of the ionizing luminosity produced may be coming from these very small regions. \\item The similar trends of dynamically \\,``aged\\,'' HII galaxies, and their relative positions in the [R -- $\\sigma$] diagram support the hypothesis of a possible evolutionary link between the two types of galaxy. If this is the case, dEs could be the descendants of HII galaxies. \\end{itemize}" + }, + "9701/astro-ph9701183_arXiv.txt": { + "abstract": "The Infrared Space Observatory (ISO), launched in November 1995, allows us to measure the far-infrared (far-IR) emission of quasars in greater detail and over a wider energy range than previously possible. In this paper, preliminary results in a study of the 5--200 $\\mu m$ continuum of quasars and active galaxies are presented. Comparison of the spectral energy distributions show that, if the far-IR emission from quasars is thermal emission from galaxian dust, the host galaxies of quasars must contain dust in quantities comparable to IR luminous galaxies rather than normal spiral galaxies. In the near-IR, the ISO data confirm an excess due to a warm `AGN-related' dust component, possibly from the putative molecular torus. We report detection of the high-redshift quasar, 1202-0727, in the near-IR indicating that it is unusually IR-bright compared with low-redshift quasars. ", + "introduction": "Quasars are multi-wavelength emitters, emitting roughly equal amounts of radiation throughout the whole electromagnetic spectrum from far-IR through to $\\gamma$-ray energies. 10\\% are also strong radio emitters. To understand the energy generation mechanisms at work, it is first essential to obtain multi-$\\lambda$ data covering the full spectral range of the emission. We now have a good understanding of the spectral energy distributions (SEDs) of low-redshift quasars and active galaxies. However, in the far-IR this has been limited by the short lifetime and wide beam of the ground-breaking IRAS satellite. Now, more than 10 years later, ISO is providing us with the chance for a second, more detailed look at the far-IR sky, allowing us to extend our knowledge to IR-fainter and higher redshift sources and to longer and shorter wavelengths (5--200$\\mu m$). To this end we are observing a sample of quasars and active galaxies with the photometer on ISO (ISOPHOT). The sample was originally designed to include $\\sim 130$ quasars and active galaxies covering the full range of redshift and of known SED properties. With the reduced in-flight sensitivities of ISOPHOT, our program has been reduced significantly and will likely include $\\sim 50$ objects, not all with full wavelength coverage. The sample will include full wavelength coverage for a well-defined subset of optically-selected, PG quasars (\\cite{alea}), along with a few high-redshift quasars, X-ray selected Seyfert 1 galaxies, and red quasars. One question that the ISO data will address is particularly relevant to this conference, namely the contribution of the host galaxy in the far-IR. Figure~\\ref{galsed} shows SEDs of spiral and elliptical galaxies superposed on the SED of a median low-redshift quasar (from \\cite{QED1}). The plot clearly shows the near-IR ($\\sim 1-2 \\mu m$ ``window\" on the host galaxy which has been used to great advantage (\\cite{mr95,Dunlop93}). Although the strength of the far-IR peak, due to cool dust, is as yet unknown, Figure~\\ref{galsed} demonstrates that this is the most likely wavelength range for a second ``window\" on the host galaxy. \\begin{figure}[h] \\psfig{figure=seds.eps,height=3.0truein} \\caption{The Median SED of a low-redshift quasar superposed on the SED of a spiral galaxy showing the well-explored, near-IR ``window\" ($\\sim 1 \\mu m$) and the potential far-IR ``window\" on a quasar's host galaxy (courtesy Kim McLeod).} \\vskip -0.5in \\label{galsed} \\end{figure} ", + "conclusions": "Although the current status of the ISO far-IR data limits the usefulness of ISO to study the host galaxy dust contribution, we can already demonstrate that, if the far-IR emission of bona fide quasars (L$> 10^{44}$ ergs$^{-1}$) is from the host galaxy, these galaxies are unusually far-IR bright, comparable to IR-bright galaxies or ULIRGs (L$_{IR} \\sim 10^{10-11.5}$ \\lsun). The ISO data also allow us to investigate the mid-IR ``AGN\" bump in quasars covering a range of redshift and luminosity. We plan to use these data to test and constrain current models of emission from a molecular torus (\\cite{pk92,rr95}). We have detected the z=4.69 quasar, 1202-0727, in the rest-frame near-IR at a level far above that seen in typical low-redshift quasars. Observations of more high-redshift quasars are necessary to determine whether the near-IR emission is unusual, as are many other aspects of this source. For pure host galaxy far-IR emission, the host must be comparable to an ULIRG to explain the mm data (\\cite{Iea94})." + }, + "9701/astro-ph9701076_arXiv.txt": { + "abstract": "In this contribution we discuss briefly a few calibration items relevant to the data analysis and present some preliminary scientific results. The discussion on instrumental topics focuses on the response matrix and Point Spread Function (PSF). In the scientific results section we discuss a first analysis of the two Seyferts MCG 6-30-15 and NGC 4151 and of the Cosmic X-ray Background. ", + "introduction": "During the Science Verification Phase (SVP) a number of bright X-ray sources have been observed by BeppoSAX. The aim of such observations is twofold: perform inflight calibrations of the instruments and verify the capabilities of BeppoSAX to achieve the scientific goals for which it has been designed. In this presentation we concentrate on the analysis of SVP data from the Medium Energy Concentrator Spectrometer (MECS) on board BeppoSAX. A more general presentation is given by Piro et al. in these proceedings. ", + "conclusions": "" + }, + "9701/astro-ph9701130_arXiv.txt": { + "abstract": "The most active starbursts are found in galaxies with the highest IR luminosities, with peak star formation rates and efficiencies that are over an order of magnitude higher than in normal disk systems. These systems are almost exclusively on-going mergers. In this review I explore the conditions needed for interactions to experience such a phase by comparing two systems at similar stages of merging but quite different IR luminosities: NGC 4038/9 and Arp 299. These observations show that the most intense starbursts occur at the sites with the highest gas densities, which is a general result for IR luminous mergers. Observations and theory both suggest that the strength of the merger induced starburst depends on the internal structure of the progenitors, the amount and distribution of the gas, and the violence of the interaction. In particular, interactions involving progenitors with dense bulges, gas-rich disks, and/or a retrograde spin are expected to preferentially lead to large amounts of gaseous dissipation, although the interplay between these parameters is unknown. A major outstanding question is how the effects of feedback alter these conclusions. ", + "introduction": " ", + "conclusions": "" + }, + "9701/astro-ph9701013_arXiv.txt": { + "abstract": "We analyze three UV images covering a $\\sim100$ square degree field toward the Virgo cluster, obtained by the FAUST space experiment. We detect 191 sources to a signal-to-noise ratio of 4.4 and identify 94\\% of them. Most sources have optical counterparts in existing catalogs and about half are identified as galaxies. Some sources with no listed counterpart were observed at the Wise Observatory. We present the results of low resolution visible spectrophotometry and discuss the foreground 101 stellar sources and the 76 detected galaxies, both in the cluster and in the fore- or background. We derive conclusions on star-formation properties of galaxies and on the total UV flux from discrete and diffuse sources in the cluster. We test for the presence of intra-cluster dust, determine the clustering properties of UV emitting galaxies, and derive the UV luminosity function of Virgo galaxies. ", + "introduction": "The ultraviolet (UV) segment of the spectrum is an excellent tool to identify and quantize active star formation (SF) regions (Donas \\etal 1987). It is also a sensitive probe of evolved hot stellar populations (\\eg \\, Burstein \\etal 1988). Therefore, UV observations may serve to understand processes of SF and stellar evolution. These processes are amplified and enhanced in a cluster environment, due to interactions among galaxies, and between galaxies and the intracluster medium. In the case of the Virgo cluster (VC) the observations are optimized by the proximity of the region, despite the VC not being a fully relaxed cluster. Throughout this work we {\\bf assume} a uniform distance of 18 Mpc to all the galaxies in the VC. A smaller cluster distance, $\\sim14$ Mpc, is suggested by the [OIII] distances to planetary nebulae (Jacoby \\etal 1990), but Visvanathan and Griersmith (1979) give 17.9 Mpc, in agreement with Sandage and Tammann (1974). We note that the depth of the cluster ($\\sim$4 Mpc) is significant in comparison to the distance to the cluster. The VC was studied intensively by many groups; we shall not give here a thorough review of all publications but mention in particular the extensive compilation of Binggeli, Sandage and Tammann (1985, BST). Their VC catalog relied mainly on the morphological appearance of a galaxy on large plate-scale photographs, to determine whether an object is member of the VC or lies in its fore- or background. In addition, the many papers by Hoffman and collaborators (Hoffman \\etal 1989a, 1989b; Hoffman 1989; Hoffman \\etal 1995) used HI observations from the Arecibo Observatory to determine the large-scale-structure in the direction of the VC. The combination of HI and optical redshifts for many of the galaxies indicates the existence of a void in the distribution of galaxies located behind the VC and in front of the Great Wall at 5500$\\leq$cz$\\leq$13500 km s$^{-1}$. The complex structure of the VC region was analyzed by Hoffman \\etal (1989a) in the context of the blue compact dwarf galaxies (BCDs) in the cluster. Their Fig. 4 emphasizes the positions of the various groups and clouds of galaxies in this region. They also established upper limits to the number of HI clouds devoid of optical counterparts in the cluster, and considered other observational properties of these galaxies, such as their FIR emission from IRAS and optical properties. Hoffman \\etal (1995) studied the background galaxies beyond the VC, in the context of their clustering properties. A reanalysis of the velocity distribution of VC galaxies (Binggeli, Popescu and Tammann 1993) showed a slight difference in the cluster redshift (1050$\\pm$35 km s$^{-1}$) from the previously accepted value (1094 km s$^{-1}$; Binggeli, Tammann and Sandage 1987). In addition to the void behind Virgo mentioned by Hoffman \\etal (1995), they found a significant lack of galaxies behind Virgo, from 3000 to 3500 km s$^{-1}$. We consider here all galaxies in the direction of the VC and with v$_{\\odot}<3000$ km s$^{-1}$ as cluster members, disregarding their membership in the Virgo core or the various sub-clusters. X--ray observations with the ROSAT PSPC (Trinchieri \\etal 1994) showed that at least one Virgo cluster galaxy (NGC 4636) may be interacting with VC gas to produce significant X-ray flux at $\\sim$3 Mpc from M87. Other such galaxies may be M86 (Forman \\etal 1979) and possibly NGC 4472 (Trinchieri \\etal 1986). In addition, we mention the discovery of very soft X-ray emission from M87 (Lieu \\etal 1996) by the EUVE satellite, with 31 counts per 1000 sec in the 100\\AA\\, band. The X-ray and far-infrared emission from M87, together with subtle optical asymmetries seen mainly in the blue, induced White \\etal (1991) to propose that this galaxy is stripped of its interstellar material by some hot gas, present at the core of the VC. High resolution HST observations of \\Lya \\, absorptions in the UV spectrum of 3C273 at the redshift of the VC (Weymann \\etal 1995 and references therein) indicate that the absorptions are probably produced by low-column-density clouds with low metallicity. One need they clearly identify is to evaluate the metagalactic background ionizing radiation field, away from the vicinity of normal galaxies; we attempt this by using new UV observations from approximately the entire Virgo cluster, obtained by the FAUST payload. \\subsection{Previous UV observations} A few galaxies in the Virgo cluster were observed in the vacuum UV by the ANS satellite (Wesselius \\etal 1982). More were observed by the rocket payloads of Smith and Cornett (1982) and Kodaira \\etal (1990), and recently by the FAUST imager (see section 2, below) on board the ATLAS--1 Shuttle experiment (Bowyer \\etal 1993). Additional measurements from a balloon--mounted telescope were reported by Donas \\etal (1987). Smith and Cornett (1982) observed the cluster with a wide field camera (11$^{\\circ}.4$ FOV) with a CsTe intensifier coupled with photographic film behind a quartz window, which define a passband with FWHM$\\simeq$1150\\AA\\, and constant--energy mean wavelength of 2421\\AA\\,. Their absolute calibration relies of $\\rho$ Vir, calibrated during the TD--1 mission against Vega, and UV magnitudes or upper limits are given for 201 galaxies. Donas \\etal (1987) used the SCAP 2000 balloon--borne telescope to image various regions of the sky in a band $\\sim$125\\AA\\, wide centered near 2000\\AA\\,. The SCAP telescope was a 17 cm diameter reflector with a 6$^{\\circ}$ field-of-view, imaging on film behind a UV-to-visible image converter/amplifier. Among the galaxies listed by Donas \\etal 12 are members of the Virgo cluster. The GUV experiment (Onaka \\etal 1989) consisted of two 17 cm diameter imagers with a 4$^{\\circ}$ field of view and was launched in 1987 on board a sounding rocket. The GUV sensitivity response had a FWHM$\\simeq$200\\AA\\, and peaked near 1450\\AA\\,. The results relevant to Virgo cluster galaxies were presented by Kodaira \\etal (1990) and a discussion of the diffuse UV light (mainly galactic) in the direction of the Virgo cluster was given by Onaka and Kodaira (1991). The galaxies' paper gives UV fluxes for 42 galaxies observed by GUV and upper limits for another 47 galaxies. Note that the reported final angular resolution of the GUV experiment was 16'.0$\\times$8'.2 (Kodaira \\etal 1990) and the equivalent exposure time on the Virgo cluster was 176 sec. ", + "conclusions": "We presented a comprehensive study of three UV FAUST images in the direction of the Virgo cluster of galaxies. We detected 191 UV sources and identified them, through comparison with catalogued optical sources, with (mostly) early-type stars or galaxies. We showed that, within the completeness limit of FAUST observations, the UV stellar distribution follows that predicted by our UV sky model. Am stars were shown to be UV--faint, in comparison to normal stars of the same spectral sub-class. For galaxies, we confirmed that significant UV flux is emitted outside the central regions. We confirmed earlier results that elliptical galaxies have widely varying UV brightness, making their use in studies of high redshift galaxy samples problematic. Star formation properties were demonstrated through UV--IR and UV--HI correlations. These confirm essentially the accepted pattern of star formation. We estimated the total ionizing flux by galaxies in the VC and showed that it is consistent with earlier estimates. We could not put significant new limits on dust in the VC but our results are consistent with a ``no dust'' case. The UV galaxies were found to concentrate in the central region of Virgo. The UV luminosity function was fitted with a Gaussian. A better determination of the UV luminosity function in the Virgo cluster requires deeper UV observations." + }, + "9701/astro-ph9701155_arXiv.txt": { + "abstract": "We present results of a new study of peculiar motions of 7 clusters in the Perseus--Pisces (PP) region, using the Fundamental Plane as a distance indicator. The sample is calibrated by reference to 9 additional clusters with data from the literature. Careful attention is paid to the matching of spectroscopic and photometric data from several sources. For six clusters in the PP supercluster no significant peculiar motions are detected. For these clusters we derive a bulk motion of 60$\\pm$220 km\\,s$^{-1}$, in the CMB frame, directed towards the Local Group. This non-detection is in marginal conflict with previous Tully-Fisher studies. Two clusters in the background of the supercluster exhibit significant negative peculiar velocities, characteristic of backside infall into PP. A bulk-flow fit to all 16 clusters reveals a statistically insignificant motion of 430$\\pm$280 km\\,s$^{-1}$ towards $l=265^{\\circ}, b=26^{\\circ}$ (CMB frame). Comparison with the velocity field predicted from the IRAS 1.2Jy survey yields $\\beta = 1.0 \\pm 0.5$. We find no evidence for residual bulk motions generated by mass concentrations beyond the limiting depth of the IRAS density field. ", + "introduction": "The Perseus--Pisces (PP) supercluster, at $cz \\sim$5000\\,km\\,s$^{-1}$, lies directly opposite the apex of the large-scale streaming detected by \\cite{lb88}. If, as in the flow model of \\cite{fb88}, the local dynamics are dominated by a single `Great Attractor' (GA) beyond Hydra--Centaurus, then the peculiar velocity at PP is predicted to be $\\la 150$\\,km\\,s$^{-1}$. If however, very distant sources are responsible for large-scale motions, then the PP region would be expected to share in a bulk streaming velocity of $\\sim 500$\\,km\\,s$^{-1}$. Measurements of bulk motion in PP have been made by \\cite{w91} and by \\cite{ha:mo}. These studies, employing the Tully--Fisher (TF) relation for spiral galaxies, both argue for a large bulk motion ($\\sim 400$\\,km\\,s$^{-1}$). \\cite{cf93} support this result, and invoke very large-scale density fluctuations to account for the large coherence length of the flow. Motions in the PP region have not been well-studied using elliptical galaxy distance indicators. Here, we describe a programme to measure peculiar motions for 7 clusters in PP, using the Fundamental Plane (FP) relation for early-type galaxies. Spectroscopic and photometric data and methods are presented in full by \\cite{pp-i}. \\cite{pp-ii} present a detailed account of the FP fits and velocity field analyses. ", + "conclusions": "\\label{conclusions} We have completed a new survey of cluster motions in the PP region, using the FP as a distance indicator. Careful attention has been paid to the construction of a homogeneous, merged dataset, and to accounting for systematic uncertainties in this procedure. We derive an insignificant bulk motion (--60$\\pm$220 km\\,s$^{-1}$ towards the Local Group) for 6 clusters in the main ridge of the supercluster. This result is in marginal conflict with Tully--Fisher surveys by \\cite{w91}, \\cite{ha:mo} and \\cite{cf93}. Two clusters in the background of PP show evidence for `backside infall' into the supercluster. Comparing the observed cluster motions with the velocity field predicted from the IRAS 1.2Jy survey, we find a best fit for $\\beta = 1.0\\pm 0.5$. There is no evidence for residual bulk flows generated by sources beyond the IRAS density field limit. A new, all-sky, survey of $\\sim$\\,50 Abell clusters within 12000\\,km\\,s$^{-1}$ is currently in progress (see Smith et al. 1997b, this volume). This improved sample, will yield useful constraints on $\\beta$, in addition to a reliable measurent of the bulk motion." + }, + "9701/astro-ph9701069_arXiv.txt": { + "abstract": "Mergers of globular clusters (GCs) once associated with dwarf spheroidal (dSph) galaxies have recently been suggested as an explanation for the bimodal horizontal branches (HBs) of some Galactic GCs, most notably NGC 1851, NGC 2808, and NGC 6229. Through analysis of the available color-magnitude diagrams for the GCs in the Fornax and Sagittarius dSph satellites of the Milky Way, as well as their metallicity distributions, we argue that the merger of two GCs would most likely produce a bimodal distribution in red giant branch (RGB) colors, or at least a significant broadening of the RGB, due to the expected difference in metallicity between the two merging globulars. No GC with a bimodal RGB is currently known, and the tightness of the RGB sequences in the above bimodal-HB GCs implies that a merger origin for their HB bimodality is unlikely. ", + "introduction": "The second-parameter phenomenon is a very intriguing anomaly affecting the color-magnitude diagrams (CMDs) of Galactic globular clusters (GCs). The morphology of the horizontal branch (HB) is observed to be primarily a function of the metallicity ${\\rm [Fe/H]} = \\log\\, ({\\rm Fe/H})_{\\rm GC} - \\log\\, ({\\rm Fe/H})_{\\sun}$, the ``first parameter.\" However, superimposed on this main trend one finds intrinsic scatter, caused by cluster-to-cluster variations in some unknown ``second parameter\" (Sandage \\& Wallerstein 1960; Sandage \\& Wildey 1967; van den Bergh 1967). The idea that age is the second parameter has been very popular, but evidence has been mounting suggesting a much more complex scheme where several parameters may be simultaneously playing more or less important r\\^oles (see Stetson, VandenBerg, \\& Bolte 1996 and Fusi Pecci \\& Bellazzini 1997 for recent reviews). One of the strongest arguments against age as the ``only\" second parameter is provided by the remarkable existence of the second-parameter effect within {\\em individual} GCs (Rood et al. 1993), such as NGC 1851, NGC 2808, and NGC 6229 (Walker 1992b; Ferraro et al. 1990; Borissova et al. 1997). Since all stars within any individual GC presumably have the same age, a different second parameter must be responsible for the existence of both red and blue HBs within the same cluster. In a recent paper, van den Bergh (1996, hereafter VDB96) has tentatively proposed that this argument against age as the second parameter may be incorrect. In his new scenario, the merger between two GCs of different HB morphologies within dwarf spheroidal galaxies (dSph's) such as Fornax or Sagittarius could lead to bimodal HB types without violating the idea that age is the second parameter. Thus, the bimodal-HB clusters could have originated in mergers between second-parameter pairs of GCs in dSph's that were subsequently accreted by our Galaxy, as envisaged by Searle \\& Zinn (1978). As noted by VDB96, the merger of two GCs could affect not only the HB morphology but also other aspects of the CMD. The purpose of this {\\em Letter} is to examine the implications of the merger scenario for the morphology of the CMD, especially in the red giant branch (RGB) region. We begin in the following section by presenting CMDs produced by ``merging\" GCs in the Sagittarius and Fornax dSph's and in the LMC. We argue that these merged CMDs are not consistent with the tight RGBs observed in Galactic GCs. In Sect. 3 we discuss the metallicity dispersions that one would expect within merged clusters. Finally, in Sect. 4 we summarize our results. ", + "conclusions": "The above analysis indicates that, in all likelihood, mergers of GCs cannot have been responsible for the HB bimodality in NGC 1851, NGC 2808, and NGC 6229, since the size of the metallicity dispersion within these clusters is very tightly constrained, as opposed to what would be expected in the merger scenario. Moreover, bimodal RGBs and RGBs which are substantially wider than those which are usually observed would be much more frequent in the Galactic GC system as a whole, if mergers of bona-fide old GCs were responsible for the existence of bimodal-HB clusters. We suggest that signs of merger events between old GCs should be searched for primarily in the RGB region of the CMD, where even small metallicity differences between the merging globulars are expected to greatly affect the resulting distributions." + }, + "9701/hep-ph9701240_arXiv.txt": { + "abstract": " ", + "introduction": "It has been known for several years that many cosmic string models appear to exhibit superconductivity~\\cite{Witten}. The currents may be very large for strings formed during the Grand-Unified phase transition, and may have significant cosmological implications. For instance, it has been argued that decaying superconducting loops can provide an explosive mechanism for galaxy formation \\cite{Ostriker} (although it is not clear if this scenario is compatible with the observed anisotropies in the microwave background \\cite{Levin}). Observational bounds on cosmic rays from superconducting strings networks \\cite{Hill} and the synchroton radiation signatures of these topological defects have been also discussed \\cite{Chud}. String loops stabilised against collapse by the effect of a current flowing on them -- {\\it vortons} \\cite{vortons} -- are a potential dark matter candidate, and indeed if they exist from the time of the GUT scale they would present an overdensity problem that could rule out string scenarios \\cite{Everyone}. In his original paper \\cite{Witten}, Witten proposed two mechanisms for cosmic string superconductivity. The first is based on a charged {\\it scalar boson} field which condenses in the core of the string, breaking electromagnetic gauge invariance there and supporting supercurrents in a manner directly analogous to conventional BCS theory. The second involves charged {\\it fermionic} zero modes trapped in the core of the string and travelling along it at the speed of light. Both mechanisms are shown to operate in a simple U(1)$\\times$U(1) model, although for the scalar-bosonic case the presence of superconductivity requires a rather intricate fine-tuning of parameters -- indicating that this is unlikely to be a generic feature of string-forming models. Despite this, as the model is sufficiently simple it has been studied thoroughly and the astrophysical consequences of superconducting strings have been explored by many authors. We refer the reader to the reviews \\cite{reviews} for more details. A distinct and more natural mechanism for string conductivity arises in non-Abelian string models, that is free of the coupling dependence that restricts the mechanisms above. Here the charge carriers are (super-heavy) vector bosons which, if absent in the model mentioned above, are present in any Grand-Unified model. This possibility was first mentioned by Preskill \\cite{Presk} and it was subsequently considered to some extent by Everett \\cite{Everett} and Alford {\\it et al.} \\cite{Alford}. By vector-boson conductivity we are {\\it not} referring to the phenomenon of W-condensation near an ordinary GUT cosmic string, which could be the source of electromagnetic currents at the electroweak scale (see, for instance, references \\cite{Amb}-\\cite{AnneWarren2}), but rather the case in which the string itself is constructed from electromagnetically charged vector bosonic fields. In this paper we will attempt a more systematic study of the conduction properties of this type of string. In section \\ref{section:Witten} we will briefly review the basics of scalar boson superconductivity in the U(1)$\\times$ U(1) model, analysing the cases of a straight infinite string and a string loop. This provides the motivation for some distinctions made later in the paper. In section \\ref{section:simple_z2_string} we introduce a toy model (SU(2)$\\times$U(1) with a complex Higgs in the adjoint representation) which is the simplest theory that supports conducting non-Abelian strings, and allows us to study their electromagnetic properies without going through the technical complications of examining a realistic Grand-Unified group. We show in section \\ref{section:SO_10}, however, that this model can be embedded directly in SO(10) with the Higgs in the {\\bf 126} representation, which is the simplest Grand-Unified theory exhibiting topologically stable strings. Finally, in section \\ref{section:conclusion} we explain why the electrodynamic properties of loops of non-Abelian conducting string appear to be very different to those of their Abelian counterparts, speculate about the cosmological implications and summarize our results. ", + "conclusions": "\\label{section:conclusion} We have seen that the essential feature of non-Abelian string conductivity models is that the string generator does not commute with electromagnetic charge. If it did, the Higgs would be annihilated by the charge generator at all distances from the string core, and so would be uncharged. It is precisely because $T_s$ and $Q$ do not share a common set of eigenvectors, and that different eigenvectors of $T_s$ have different profiles (in the discussion in section \\ref{section:simple_z2_string}, $f$ and $h$) near the string core, that a charged component of the Higgs exists there. Charge is, however, generically ill-defined at the origin. The charged component of the Higgs (see equation \\ref{eqn:base_decomp}) cannot really be considered to be a well-defined `condensate' in the sense of the Witten model. Moreover, we cannot universally give this charged component a varying phase with an electromagnetic gauge transformation around a loop of such string, as it does not vanish at $r=0$ where $Q$ itself is ill-defined. If we were to try to construct such a solution we would have to give the Higgs a phase with a generator that commuted with $T_s$ at $r=0$ and deformed into $Q$ outside the string. That implies there must be some radial dependence of the Lie algebra element that acts on the condensate, and we cannot construct a non-trivial $r$ and $z$-dependent condensate that is periodic in $z$, as is required for a loop. This can be seen as follows. Following the infinite-string ansatz of Alford {\\it et al.} \\cite{Alford} for a zero-mode excitation, we require a Higgs of the form \\[ \\Phi(r,\\theta, z,t) = e^{i\\beta(z,t){\\bf S}(r,\\theta)}\\bar{\\Phi}(r,\\theta), \\] where $\\bar{\\Phi}$ is the unperturbed string solution, and ${\\bf S}$ lies in the Lie alebra of the unbroken group $H$ and tends to charge $\\bf{Q}(\\theta)$ away from the string core. We wish to implement this on the topology of a loop of length $L$. Now the coordinate $z$ measures the distance along the string. The requirement \\[ \\Phi(z+L) = \\Phi(z) \\] implies that \\be \\beta(z+L,t){\\bf S}(r,\\theta) \\equiv \\beta(z,t){\\bf S}(r,\\theta) + 2\\pi N \\mbox{\\boldmath $1$}. \\ee This cannot be satisfied everywhere unless $\\beta$ is trivial, or $S$ is independent of $r$ and $\\theta$ (which is not possible for the case $[T_s,Q] \\neq 0$ we require). We see, then, that we must view current excitations on such non-Abelian strings as being supported by the string gauge field, which is also charged, and sources excitations of the form \\be {\\bf A}_\\alpha = -{\\bf S}(r,\\theta) \\beta (z,t). \\ee This viewpoint is consistent with the topology of a string loop. The boundary conditions on the gauge field equations of motion (which, in the absence of a Higgs phase, are homogeneous in $S$) now demand that the ${\\bf A}_\\alpha$ vanish both at the center of the string loop and at spatial infinity \\cite{Landau}. The lowest energy solution in the absence of an external field is simply ${\\bf S}\\equiv 0$ and hence we suggest that an uncharged non-Abelian string loop ($\\epsilon > 0$) will be incapable of supporting a persistent current in this case, making it in effect a perfect conductor. We aim to explore the issue of perfect conductor electrodynamics in a subsequent paper\\footnote{Work in progress.}. Similarly, we would not expect the charged $\\epsilon < 0$ case to persist either. It is conceivable that the chiral state, in which the current is identified with the charge density, {\\it could} be dynamically stable, however. The solution of the equations of motion on a string loop need to be solved fully to confirm this, however, to take into account curvature effects. Studies of current-carrying loop stability have to date worked with the Witten model of a superconducting string, in which a loop (or {\\it vorton}) has two associated conserved quantities $N$ and $Z$; these correspond to the condensate winding and the net charge carried by the loop. In terms of the function $\\beta(z,t)=kz \\pm \\omega t$ defined on the string worldsheet, these come from the first term (a topological invariant) and second term (Noether charge) respectively. The fact that there is no condensate in non-Abelian models seems to preclude the existence of the quantum number $N$. However, it may be that the {\\it chiral} current-carrying states may be dynamically stable on a loop, with a net charge $Q$. These may be vulnerable to the loss of charge carriers at cusps \\cite{Cop}, however, and so these currents can leak from the string when the loop has a contorted small-scale structure. Despite the questions raised about the stability of currents, it seems that there are so many competing perfect-conduction mechanisms clamouring for a string's attention that it seems almost certain that if strings exist at all they will be current-carrying at some scale. It seems that if GUT strings are formed they are likely to be capable of carrying enormous currents -- we see that the Euler-Lagrange equations provide no natural upper limit, and so quantum effects must intervene through particle production in the vicinity of the string. Chiral current magnitudes may be limited by the amount of charge a loop can randomly acquire, either at formation or after intercommutation of strings carrying different currents. If the bare GUT string does not support currents, whether fermionic or vector bosonic, then subsequent dressing at the electroweak scale appears to provide them. This provides a further degree of freedom for string-based models of structure formation and dark matter. Note that the {\\it induced} vector-boson conductivity models \\cite{AnneWarren1}-\\cite{AnneWarren2} are not included in this. They are complementary to the models listed and will presumably display similar topological currents and/or zero modes according to whether charge commutes with the string generator or not. In this paper we have explored non-Abelian string conductivity in some detail. We have emphasised the role of topology in true superconductivity, and its absence in non-Abelian string models. We have analysed the current solution (with and without backreaction) in an illustrative non-Abelian string model. We have also shown that the standard, lowest-energy SO(10) string solution is a simple embedding of this model and hence also supports currents. We suggest that only charged non-Abelian strings will be able to carry persistent currents in the absence of an external field, due to dynamical stability of the chiral zero-modes. It seems that vorton stability is less likely in realistic GUT string models, then, which avoids the overdensity problem and reinforces structure formation scenarios involving defects at such a scale." + }, + "9701/astro-ph9701188_arXiv.txt": { + "abstract": "The Supernova Cosmology Project has discovered over twenty-eight supernovae (SNe) at $0.35 0.3$) Hubble constant, $H_0^G$ (Bartlett et al. 1995).\\nocite{ba:h30} Alternatively, it may be that these Hubble constant measurements lie on the tail of their statistical and systematic error distributions. We use our first sample of seven $z>0.35$ Type Ia supernovae (SNe Ia) to address both these possibilities, first by directly comparing our SN Ia sample with one lying within the local Hubble flow to determine the ratio of $H_0^L$ to $H_0^G$, and then by using our sample (the first SNe observed in this redshift regime) together with SN Ia absolute magnitude calibrations to determine the value of $H_0^G$. The possibility that $H_0^L/H_0^G \\ne 1$ has arisen in the context of the observation of peculiar velocity fields (e.g., de Vaucouleurs 1958; Dressler et al. 1987; Lynden-Bell et al. 1988). Simulations of \\cite{tu:hubble} have shown that measured Hubble constants depend on the observer location and the depth of observations. Previous work by \\cite{lauer:1992} has constrained deviations from uniform Hubble flow to be $\\Delta H_0/H_0 < 0.07$ at $0.01\\le z \\le 0.05$ using brightest cluster galaxies as a distance indicator. The same sample of galaxies shows evidence for a peculiar motion of $689 {\\rm~km~s^{-1}}$ with respect to the cosmic background radiation (Lauer \\& Postman 1994),\\nocite{lauer:1994} although \\cite{ri:1995lg} argue that SNe Ia at similar redshifts do not support this conclusion. We thus must still examine the possibility of a large scale ($z \\ge 0.05$) peculiar velocity flow affecting all the local $H_0$ measurements. The Supernova Cosmology Project has discovered over twenty-eight SNe in the redshift range $0.35 < z < 0.65$ in a systematic search (Perlmutter et al. 1994; 1995).\\nocite{sne1994} \\nocite{sne1995} The peak magnitudes of these high-redshift candles, when compared with the peak magnitudes of local SNe, can yield measurements of the cosmological parameters $\\Omega_M$ and $\\Lambda$ (Goobar \\& Perlmutter 1995; Perlmutter et al. 1996b). \\nocite{goo:lambda}\\nocite{pe:1996} This calculation implicitly assumes that the local SN calibrators lie within the global cosmological flow; i.e., that we do not live in a local bubble where peculiar velocities appreciably bias the observed value of the Hubble constant. In this paper we take an alternative approach, leaving $\\Omega_M$ and $\\Lambda$ as free parameters and using our high redshift SNe Ia to measure the ratio between the locally observed Hubble constant and the global Hubble constant, $H_0^L/H_0^G$. We also use our SNe Ia to obtain a measurement of the Hubble constant. This can be compared to the other SN-based measurements which range from ${\\rm~57~km~s^{-1} Mpc^{-1}}$ (Sandage et al. 1996) \\nocite{sandage:1996} to $\\sim 66 {\\rm~km~s^{-1} Mpc^{-1}}$ (Hamuy et al. 1995; Riess, Press, \\& Kirshner 1996), and to the \\nocite{ha:hubble}\\nocite{ri:lcs2} above mentioned Cepheid methods that connect distances in a sequence from a single galaxy, to the core of its cluster, and then to the Coma cluster. We use the first seven SNe from our search. A detailed description of our search methodology, the telescopes used, the photometric and spectroscopic data compiled for each event, light curve analysis, and a study of possible systematic uncertainties, are given in Perlmutter et al. (1996a,b).\\nocite{pe:1996} \\nocite{pe:aig} Specifically, we use the redshift as measured from the host galaxy spectrum, the best fit K-corrected $B$ peak magnitude after our galaxy extinction correction $m_B=m_R-K_{BR}-A_R$, the value of $\\Delta m_{15}$ (Phillips 1993)\\nocite{ph:1993}, and $m_B$ after correction to the Leibundgut template $m_B^{\\{1.1\\}}$ using the relation of \\cite{hamuy:1996} as discussed in \\S 2. ", + "conclusions": "The measurement of cosmological distances using high-redshift SNe with locally-calibrated standard candles sets a limit on the differences between the local and global Hubble constants. From our analysis, it is clear that these data are inconsistent with scenarios that use a local bubble with high $H_0^L$ that differs greatly from $H_0^G$. We also obtain an upper limit for the Hubble constant that is consistent with many of the other current measurements. However, limits that disagree with higher $H_0^G$ measurements may be obtained with independent upper limits on $\\Omega_\\Lambda$. The SN Ia absolute magnitude calibrations are still subject to debate and may have systematic errors larger than the statistical ones given above, so it is important to ask how robust our results are. An uncertainty in the absolute calibration $\\delta m$ in magnitudes propagates into $\\delta H_0/H_0 \\approx \\delta m$. A 0.09 mag difference in the magnitude calibrations, such as the one between the $\\Delta m_{15}$-corrected absolute magnitudes for six SNe, $M_B^{\\{1.1\\}}=-19.45 \\pm 0.07 {\\rm~mag}$, and that of \\cite{ri:lcs2}, $M_{B,\\Delta=0}=-19.36 \\pm 0.1 {\\rm~mag}$, will produce a $10 \\%$ change in either $H_0^G$ or $H_0^L/H_0^G$. There is little difference between magnitude corrected and uncorrected results for the ratio $H_0^L/H_0^G$, but there is a systematic difference for $H_0^G$ itself, as seen in Table~\\ref{bounds}. This is because both the light-curve-width distribution and the width-magnitude relation of our high-redshift sample are similar to the distribution and relation of the \\cite{hamuy:1996} sample but not to those of the \\cite{sandage:1996} sample. Although these differences may be due to selection effects, the small number statistics of the Cepheid-calibrated SN sample can also produce fluctuations that account for the differences. In \\cite{pe:1996} we calculated $\\Omega_M$ and $\\Omega_\\Lambda$ setting $H_0^L$ equal to $H_0^G$, whereas in this paper we have discussed the measurement of $H_0^L/H_0^G$ while leaving $\\Omega_M$ and $\\Omega_\\Lambda$ as free parameters. Ideally one would like to measure both sets of quantities simultaneously. (This problem has been discussed in Wu, Qin, \\& Fang 1996.)\\nocite{wu:1996} Filling in a Hubble diagram with measurements of spatially well-distributed SNe should make it possible to decouple local and global streaming motions by showing redshift dependent deviations from the standard model, and allow one to measure $\\Omega_M$ and $\\Omega_\\Lambda$ independently of local peculiar flows. Using SNe from redshift regimes with no evidence of flows, we can simultaneously fit $H_0^G$, $\\Omega_M$, and $\\Omega_\\Lambda$ using Equation~\\ref{answer2}, thus producing a measurement of the Hubble constant. Our current data set, which spans from $0.35 < z < 0.5$, shows no sign of peculiar flows but needs higher statistics and more complete spatial coverage to confirm this result. This work was supported in part by the National Science Foundation (ADT-88909616, AST-9417213) and the U.~S. Department of Energy (DE-AC03-76SF000098). \\clearpage" + }, + "9701/astro-ph9701141_arXiv.txt": { + "abstract": "We present new age and distance determinations for the Galactic Globular Clusters M55 and M5, using the luminosity function method (Jimenez \\& Padoan 1996, Padoan \\& Jimenez 1997). We find an age of $11.8 \\pm 1.5$ Gyr for M55 and $11.1 \\pm 0.7$ Gyr for M5. This confirms previous results (Jimenez et al. 1996, Sandquist et al. 1996) and allows to conclude that {\\bf the oldest stars in the Universe are not older than 14 Gyr}. We also find $m-M=14.13 \\pm 0.11$ for M55, and $m-M=14.49 \\pm 0.06$ for M5. These values agree with the ones obtained using the tip of the red giant branch (Jimenez et al. 1996) and the sub-dwarf fitting method (Sandquist et al. 1996). ", + "introduction": "An accurate determination of the ages and distances of Globular Clusters (GCs) is an important constraint for the age of the Universe, and for the theory of galaxy formation. In particular it is important to compute very accurate relative ages to understand if there is a spread in ages among the Galactic GCs or not. The use of the stellar luminosity function (LF) to compute ages of GCs was first proposed by Paczynski (1984). Later on, Jimenez \\& Padoan (1996) and Padoan \\& Jimenez (1996) developed a method to determine {\\it the age and the distance of a GC simultaneously}, using the LF. The method is described in detail in Padoan \\& Jimenez (1996), where it is concluded, on the basis of artificial data, that an uncertainty of about 0.6 Gyr in the age and 0.06 mag in the distance modulus can be achieved, if the number of stars, in 1 mag-wide luminosity bins, is known with an uncertainty of 3\\%. In this paper we use recent observations of the Galactic Globular Clusters M5 (Sandquist et al. 1996) and M55 (Desidera \\& Ortolani, private communication) to apply the LF method and compute accurate ages and distance module. These two clusters are very adequate since M55 is a metal-poor one and M5 has intermediate metallicity, so we can investigate the spread in ages (if any) in the formation of the GC system. In this letter we apply for the first time the LF method to real data and we show that the method is much more superior to traditional methods (isochrone fitting to the main sequence turn off point, $\\Delta V$ method, or any other methods that involve the fitting of the main sequence turn off). The method is superior because it allows to determine the age and the distance simultaneously and independently and because the errors in computing the age and distance are straightforward to calculate. Furthermore, it gives age determinations with sufficient accuracy to make cosmological predictions. This first application of our LF method to real data shows that our previous theoretical predictions were correct. ", + "conclusions": "The results listed in Table~1 show that the age and the distance modulus of M55 are in good agreement with previous determinations by Mandushev et al. (1996) and Alcaino et al. (1992). \\begin{table} \\begin{center} \\begin{tabular}{ccc} & M5 & M55 \\\\ \\hline\\hline age & $ 11.1 \\pm 0.7$ & $11.8 \\pm 1.5$ \\\\ m-M & $ 14.49 \\pm 0.06$ & $14.13 \\pm 0.11$ \\\\ \\hline \\end{tabular} \\caption{The table gives the values for the age and distance modulus for M5 and M55. These values have been determined {\\it simultaneously} using the luminosity function method described in the text.} \\end{center} \\end{table} In the case of M5 the results agree with Sanquist et al. conclusions. They estimate in fact an age of $13.5 \\pm1$ Gyr, for $[Fe/H]=-1.17$, and they state that the age would be 11.5 Gyr, for $[Fe/H]=-1.4$. We use $[Fe/H]=-1.3$, and get an age of $11.1\\pm0.7$ Gyr. They also estimate $m-M=14.50\\pm0.07$ mag for $[Fe/H]=-1.17$, and $m-M=14.41\\pm0.07$ mag, for $[Fe/H]=-1.4$, using the sub-dwarf fitting of the main sequence. We get $m-M=14.49\\pm0.06$ mag, for $[Fe/H]=-1.3$. Note that in Padoan \\& Jimenez (1996) we estimated a variation of 0.02 mag in $m-M$, for a shift of 0.1 in metallicity; so we would predict $m-M=14.47\\pm0.06$ mag, for $[Fe/H]=-1.4$. \\begin{figure} \\centering \\leavevmode \\epsfxsize=1.0 \\columnwidth \\epsfbox{fig1.eps} \\caption[]{The figure shows the contour plots of $R(t,m-M)$ (see text) in determining simultaneously the distance modulus and age of M55. Notice that the contours closed around a central value, showing that the method works quite well in breaking the age-distance degeneracy.} \\end{figure} The distance modulus measured with the LF method is therefore in excellent agreement with the distance modulus determined with the sub-dwarf fitting. The uncertainty of our estimates is very small ($\\pm0.06$ mag), and no assumption on the age is required. The LF method is superior to the main sequence turn-off (MSTO) method (Chaboyer, Demarque \\& Sarajedini 1996), to determine the absolute age of globular clusters, because it is not affected by the three largest sources of theoretical uncertainty affecting the MSTO method, that is to say the determination of the value of the mixing length parameter, the morphology of the MSTO and the color-$T_{\\rm eff}$ calibration (see Jimenez et al. 1996 for a detailed discussion of the main uncertainties in the MSTO method). Furthermore, the MSTO method needs to know the distance in order to determine the age, and it is unable to break this degeneracy. The absolute ages, determined in this work for M55 and M5, seem to indicate that the oldest GCs are not older than 14 Gyr. The LF method is a very powerful tool to investigate relative ages, since most uncertainties of stellar evolution theories are in that case avoided. From the comparison of the ages of M5 and M55 we can conclude that the age of the two GCs is not significantly different. We conclude by remarking that most methods to determine age and distance module of GCs share two common problems: some degree of dependence of age on distance modulus (or vice-versa), and a somewhat fuzzy procedure to estimate the uncertainty of the final result. Our LF method, instead, gives constraints for both age and distance modulus independently, and estimates both most probable values and uncertainties in a straightforward way. \\begin{figure} \\centering \\leavevmode \\epsfxsize=1.0 \\columnwidth \\epsfbox{fig2.eps} \\caption[]{The same as before but for M5. The estimated uncertainty of 4\\% is marked with dashed lines. } \\end{figure} On the basis of the present work, we think that very high quality data for GCs, together with the LF method, may shed new light on the problems of the age of the oldest stars in the Universe and the formation of the Galaxy." + }, + "9701/astro-ph9701007_arXiv.txt": { + "abstract": "We use network calculations of r-process nucleosynthesis to explore the origin of the peak in the solar r-process abundance distribution near nuclear mass number $A\\approx 160$. The peak is due to a subtle interplay of nuclear deformation and beta decay, and forms not in the steady phase of the r-process, but only just prior to freezeout, as the free neutrons rapidly disappear. Its existence should therefore help constrain the conditions under which the r-process occurs and freezes out. ", + "introduction": " ", + "conclusions": "" + }, + "9701/astro-ph9701160_arXiv.txt": { + "abstract": "Recent measurements of Si~IV/C~IV ratios in the high-redshift Ly$\\alpha$ forest (\\markcite{SC96}Songaila \\& Cowie, AJ, 112, 335 [1996a]; \\markcite{S97}Savaglio et al., A\\&A, in press [1997]) have opened a new window on chemical enrichment and the first generations of stars. However, the derivation of accurate Si/C abundances requires reliable ionization corrections, which are strongly dependent on the spectral shape of the metagalactic ionizing background and on the ``local effects'' of hot stars in nearby galaxies. Recent models have assumed power-law quasar ionizing backgrounds plus a decrement at 4 Ryd to account for He~II attenuation in intervening clouds. However, we show that realistic ionizing backgrounds based on cosmological radiative transfer models produce more complex ionizing spectra between 1--5 Ryd that are critical to interpreting ions of Si and C. We also make a preliminary investigation of the effects of He~II ionization front non-overlap. Because the attenuation and re-emission by intervening clouds enhance Si~IV relative to C~IV, the observed high Si~IV/C~IV ratios do {\\it not} require an unrealistic Si overproduction [Si/C $\\geq 3$ (Si/C)$_{\\odot}$]. If the ionizing spectrum is dominated by ``local effects'' from massive stars, even larger Si~IV/C~IV ratios are possible. However, unless stellar radiation dominates quasars by more than a factor of 10, we confirm the evidence for some Si overproduction by massive stars; values Si/C $\\approx 2$ (Si/C)$_{\\odot}$ fit the measurements better than solar abundances. Ultimately, an adequate interpretation of the ratios of C~IV, Si~IV, and C~II may require hot, collisionally ionized gas in a multiphase medium. ", + "introduction": "The detection of metals associated with the ``high'' column density [N(H~I) $> 10^{14.5}$ cm$^{-2}$] Ly$\\alpha$ forest of absorbers in quasar spectra (\\markcite{C95}Cowie et al. 1995; \\markcite{T95}Tytler et al. 1995) provides new challenges to the understanding of the nature and origin of these absorbers. The contamination by metals must now be accounted for in models for cloud evolution. It is also possible to perform a detailed analysis of the thermodynamic and ionization state of these absorbers, which have previously been probed only by their hydrogen Ly$\\alpha$ absorption and the integrated effect of their He~II Ly$\\alpha$ absorption (\\markcite{GFS95}Giroux, Fardal, \\& Shull 1995). If the clouds are photoionized, the detection of C~II, C~IV and Si~IV may constrain the shape of the metagalactic radiation spectrum and indicate the epoch of dramatic changes in the ionizing radiation field at $z > 3$ (\\markcite{SC96}Songaila \\& Cowie 1996a; \\markcite{S97}Savaglio et al. 1997). The multiple C~IV lines associated with a single H~I absorption complex \\markcite{SC96}(Songaila \\& Cowie 1996a) belies a simple picture of a homogeneous cloud. These absorbers, like Lyman Limit systems, may better be interpreted with multiphase models (\\markcite{GSS94}Giroux, Sutherland, \\& Shull 1994; \\markcite{PRR96}Petitjean, Riediger, \\& Rauch 1996; \\markcite{HRS97}Haehnelt, Rauch, \\& Steinmetz 1997; \\markcite{H97}Hellsten et al. 1997) that may include hot, collisionally ionized gas. We defer such interpretation to a later study, focussing here on the tradition of analyzing metal-line ratios with single-phase photoionization models. In this way, we can isolate the effects of the incident spectrum on species with ionization potentials in the range 1--5 Rydbergs (e.g., C~III, C~IV, Si~III, Si~IV). For example, \\markcite{SC96}Songaila \\& Cowie (1996a) argue that the observed increase in the Si~IV/C~IV ratio at $z > 3.1$ suggests that the He~III regions around ionizing sources may not have overlapped by that epoch, cutting off all radiation with $h \\nu > 54.4$ eV. Here, we primarily consider models for the averaged metagalactic background in which He~II ionization fronts have overlapped before $z \\approx 4$. We do consider a limiting case in which no photons above 4 Ryd are present in the background, as well as a case where we restore the much less heavily attenuated x-ray background. In the past decade, metal-line ratios in Lyman Limit ($N_{HI} {}^>_\\sim 10^{17}$ cm$^{-2}$) absorbers have been used to probe the shape of the ionizing spectrum, and to estimate the intrinsic metallicity of the absorbers (cf. \\markcite{BS86}Bergeron \\& Stasinska 1986; \\markcite{SBS88}Sargent, Boksenberg \\& Steidel 1988; \\markcite{SSB88}Steidel, Sargent \\&, Boksenberg 1988; \\markcite{SS89}Steidel \\& Sargent 1989; \\markcite{BI90}Bergeron \\& Ikeuchi 1990; \\markcite{MO90}Miralda-Escud\\'e \\& Ostriker 1990; \\markcite{DS91}Donahue \\& Shull 1991; \\markcite{M92}Madau 1992). With the advent of new spectrographs on large telescopes, the lower $N_{HI}$ absorbers can be exploited in the same way. These absorbers hold a special interest, because the source of their metal enrichment is still unclear. While the metal lines in Lyman limit absorbers have been associated with the haloes of bright galaxies, the number of Ly$\\alpha$ forest absorbers greatly exceeds the corresponding number of bright galaxies, and the Ly$\\alpha$ absorbers do not seem to show the same clustering properties as galaxies. As yet, metal lines have only been associated with ``high-column'' absorbers [N(H~I) $>10^{14.5}$ cm$^{-2}$], which numerically are a small fraction of the Ly$\\alpha$ forest. It is still possible that these clouds may represent an overlap of the high-column end of the pristine Ly$\\alpha$ forest clouds and the low-column end of the metal enriched clouds associated with galactic haloes. As a result some, if not all, of the metal-line absorbers may still be associated with the galaxies that contain the stars responsible for their metals. For example, over a Hubble time at $z = 3$, galactic outflows at $300 V_{300}$ km~s$^{-1}$ will transport heavy elements a distance \\begin{equation} d \\leq (300~{\\rm kpc}) V_{300} h^{-1} \\left[ \\frac {1+z}{4} \\right]^{-3/2} \\end{equation} for a Friedmann universe with $H_0 = (100~{\\rm km~s}^{-1})h$ and $\\Omega_0 = 1$. In practice, the period of heavy element injection may have lasted only $10^8$ yrs, and the metals could move distances of only 30-50 kpc from their sources (bright galaxies, dwarf galaxies, globular clusters). As \\markcite{MS96}Madau \\& Shull (1996) have emphasized, metal enrichment of the Ly$\\alpha$ absorbers implies a substantial Lyman continuum (LyC) emission accompanying the star formation at $z {}^>_\\sim 3.5$. Thus, a large fraction of the metagalactic background at high redshift could be due to massive stars, a point also made by \\markcite{GS96}Giroux \\& Shapiro (1996) and \\markcite{S97}Savaglio et al. (1997). Furthermore, the absorbers with metal lines may be sufficiently close to the local sources that the incident radiation field is dominated by local hot-star radiation. In contrast, if the metal enrichment of the Ly$\\alpha$ forest arises from Population III stars at $z > 10$ (\\markcite{C85}Couchman 1985; \\markcite{OG96}Ostriker \\& Gnedin 1996), the stellar ionizing radiation will be greatly attenuated by intergalactic absorption, and the metagalactic background at $z = 3$ will be dominated by the harder spectrum of quasars. Because metal-line systems would then be less likely to be associated with nearby sources of stellar radiation, their absorption-line ratios would be more representative of the metagalactic ionizing background. Preliminary analyses (\\markcite{SC96}Songaila \\& Cowie 1996a; \\markcite{S97}Savaglio et al. 1997; \\markcite{H97}Haehnelt 1997) of the Ly$\\alpha$ forest clouds have assumed photoionization by simple power-law ionizing spectra. The only way these models account for filtering of the ionizing source spectrum by intervening clouds is to include a decrement at $\\nu_{HeII}$ (4 Ryd). More realistic models include several additional effects. First, when absorption due to intervening clouds is included, the background spectrum just above the H~I and He~II ionization edges shows a flatter power law than that of the sources (\\markcite{MO90}Miralda-Escud\\'e \\& Ostriker 1990; \\markcite{MM94}Madau \\& Meiksin 1994; \\markcite{GFS95}Giroux et al. 1995). Second, \\markcite{HM96}Haardt \\& Madau (1996) and \\markcite{FGS97}Fardal, Giroux, \\& Shull (1997) show, using cosmological radiative transfer, that cloud emission by recombination to the H~I and He~II ground states and He~II Ly$\\alpha$ and 2-photon radiation alters the extent of the decrement at $\\nu_{HeII}$ and changes the shape of the spectrum between 1 and 5 Rydbergs. Including a more realistic spectral shape will alter the relative populations of metal ions with edges in the 1 to 5 Ryd range even if the ionizing sources are all assumed to be quasars with intrinsic power-law ionizing spectra. In this paper, we consider the extent to which these alterations of the assumed shape of the incident radiation field affect the interpretation of the metallicity of these absorbers and the nature of the photoionizing sources. We use the data sets of \\markcite{SC96}Songaila \\& Cowie (1996a) and \\markcite{S97}Savaglio et al. (1997), which include Si~IV and C~IV lines associated with higher column density Ly$\\alpha$ forest clouds. \\markcite{SC96}Songaila \\& Cowie (1996a) provide an important subsample for which C~II lines are also observed. ", + "conclusions": " 1. {\\it Ionizing Radiation Field}.-- The metagalactic radiation field is likely to include both power-law (AGN) and stellar (hot-star) components. The Ly$\\alpha$ forest absorbers with metal lines may also experience a local radiation field from starburst galaxies within 50--100 kpc. 2. {\\it Si/C Overabundance}.-- For plausible mixtures of stellar and AGN spectra in the metagalactic background, our photoionization models produce enhanced Si~IV/C~IV ratios, consistent with high-$z$ absorbers with Si/C $\\approx 2$(Si/C)$_{\\odot}$ at low ionization (C~II/C~IV $>$ 0.1). For absorbers with C~II/C~IV $<$ 0.05, it is difficult to account for the high values of Si~IV/C~IV unless Si/C $>$ 10(Si/C)$_{\\odot}$, an unrealistically large value for massive-star nucleosynthesis. These systems may include photoionization from local stellar sources as well as hot, collisionally ionized gas. 3. {\\it Local Ionizing Sources}.-- If the radiation field incident on the absorbers is dominated by a nearby starburst galaxy, the Si~IV/C~IV ratios are further enhanced. If C~II/C~IV $>$ 0.1, no Si/C overabundance is necessary to explain Si~IV/C~IV ratios if absorbers are within about 40~kpc of a starburst galaxy {\\it and} the background is dominated by stellar sources. Even if the metagalactic background is dominated by AGNs, close proximity to a starburst galaxy may reduce the needed Si/C overabundance to a factor 1.5. When C~II/C~IV $^<_\\sim 0.05$, it remains very difficult to account for the high values of Si~IV/C~IV with photoionized models, although models with locally dominated radiation may only require Si/C enhancements of a factor of 2-3. 4. {\\it Non-Overlapping He~II I-Fronts}.-- Around quasars whose ionizing continua have steep spectral indices ($\\alpha_s > 1.84$), the He~II I-fronts will lie within the H~I I-fronts. If many absorbers lie in regions where all photons above 4 Ryd are attenuated (\\markcite{SC96}Songaila \\& Cowie 1996a), we obtain good agreement with almost all measurements of Si~IV/C~IV if Si/C is enhanced by a factor of 2--3. However, including even a small contribution of higher energy photons in the background increases the needed Si/C overabundance to an order of magnitude for absorbers with the lowest C~II/C~IV ratios. 5. {\\it Temperature Effects}.-- If the absorbers are cooler than expected for thermal equilibrium between photoelectric heating and radiative cooling, the Si~IV/C~IV ratio is increased for a given C~II/C~IV ratio. At low C~II/C~IV, this enhances Si~IV/C~IV by a factor of 5 in our photoionization models. Conversely, if the absorbers are hotter, Si~IV/C~IV is lower for a given C~II/C~IV ratio. \\vspace{0.5cm} While we have described several processes that might increase the Si~IV/C~IV ratio, an important additional constraint on these possibilities is the fact that the high ratios are preferentially found at $z > 3.1$. Recent observations (\\markcite{B97}Boksenberg 1997) challenge this interpretation, by finding high Si~IV/C~IV ratios in absorbers at $z = 2-3$. If, however, Si~IV/C~IV rises above $z > 3.1$, this argues for a time dependence to whatever process increases this ratio. It is this property, as well as the increased He~II absorption toward Q0302-003 at $z=3.28$ (\\markcite{J94}Jakobsen et al. 1994), that \\markcite{SC96}Songaila \\& Cowie (1996a) use to support their suggestion that the absorbers lie in regions where He~II I-fronts have not overlapped at $z > 3.1$. From a higher resolution HST/GHRS spectrum of Q0302-003, \\markcite{H97}Hogan, Anderson, \\& Rugers (1997) find evidence for residual transmitted flux below the He~II edge. Their 95\\% confidence upper limit, $\\tau_{HeII} \\le 3$, makes it less necessary to propose that Q0302-003 lies in a region where He~II I-fronts have not overlapped, but does not preclude the suggestion that such regions existed at $z>3.1$. If this is the case, radiation outside of He~III regions would largely limit the range in ionization stages for carbon to II-IV, and for silicon to II-V. In AGN, however, the large fluxes of unattenuated photons above 45 eV can ionize Si~IV and C~III, so that high observed Si~IV/C~IV ratios would be difficult to explain for large U (small C~II/C~IV). As a result, the spectrum below 4 Ryd must be dominated by stellar sources. This may be difficult to achieve with the known quasar luminosity functions at $z > 3.5$. In addition, soft x-ray radiation from quasars, which is less attenuated by intervening clouds even in the low He~III porosity case, may again make high Si~IV/C~IV, low C~II/C~IV absorbers difficult to understand if they are solely photoionized by a metagalactic background. New or improved measurements of He~II absorption may soon indicate whether He~II I-fronts have overlapped well before $z = 3.1$. In that case, a sharp change in the shape of the metagalactic radiation field may be less plausible. Still, in general, the trend is likely to be more attenuation of photons above the 4 Ryd limit with increasing redshift. Other time dependent effects that enhance the ratio of Si~IV/C~IV are possible. Higher abundances of Si relative to C are associated with metal yields from the most massive stars. However, in a flat universe with $h=0.75$, the age of the universe exceeds $10^9$ years by $z = 3.1$, time enough for lower mass stars to enrich the gas with carbon. For example, if multiple supernovae eject metal-enriched gas into the IGM, the additional $4 \\times 10^{8}$ years between $z=3.5$ and $z=2.5$ may make it more likely that the gas surrounding the high mass stars has been enriched with carbon from a previous episode of star formation. The temperature effects we discuss in \\S 4 may also have a time dependence. If the IGM has been uniformly enriched by a much earlier episode of Population III star formation, and if the Ly$\\alpha$ forest clouds with metals have their origin in growing overdensities in the IGM, then the absorbers observed at higher $z$ may have formed at an earlier epoch. If Compton cooling off the cosmic microwave background was a dominant coolant at this epoch (\\markcite{MR94}Miralda-Escud\\'e \\& Rees 1994), the cloud may have retained memory of this lower temperature. This effect is probably not large, since there is not a significant evolution in the observed line widths. Alternatively, one may resolve the puzzle of the high Si~IV/C~IV, low C~II/C~IV absorbers by relaxing the single-phase model which we used to explore the effects of radiative transfer. As we have emphasized in a previous paper (\\markcite{GSS94}Giroux et al. 1994), the production of heavy elements in QSO absorption systems is naturally accompanied by hot gas due to supernovae and hot-star winds." + }, + "9701/astro-ph9701210_arXiv.txt": { + "abstract": "\\noindent \\rightskip=0pt We calculate the anisotropies in the cosmic microwave background induced by long-wavelength primordial gravitational waves in a universe with negative spatial curvature, such as are produced in the ``open inflation'' scenario. The impact of these results on the {\\sl COBE\\/} normalization of open models is discussed. ", + "introduction": " ", + "conclusions": "" + }, + "9701/astro-ph9701032_arXiv.txt": { + "abstract": "General relativity asserts that: energy and momentum conservation laws are valid, preferred frames do not exist, and the strong equivalence principle is obeyed. In this paper recent progress in testing these important principles using millisecond pulsars is summarised. ", + "introduction": "\\def\\bb{$\\bullet$~} The fundamental physics and principles that can be observed and tested by the exceptional precision of pulsar timing includes (\\cite{bel97}): \\noindent \\bb Relativistic precession \\\\ \\bb Shapiro delay \\\\ \\bb Einstein delay \\\\ \\bb Gravitational waves \\\\ \\bb Variation in G \\\\ \\bb Chandrasekhar mass \\\\ \\bb Spin-orbit coupling \\\\ \\bb Ultra low frequency gravitational waves \\\\ \\bb Strong equivalence principle \\\\ \\bb Lorentz Invariance \\\\ \\bb Conservation laws At this meeting Esposito-Far\\`{e}se gave an update on the first 4 items and summary of the parametrised post-Newtonian formalism (PPN). Will (1993) \\nocite{wil93} also discusses the PPN formalism in detail and gives limits on many of the ten PPN parameters. Taylor et al. (1992) \\nocite{twdw92} discuss many other relativistic effects which could in principle be measured with sufficient precision. Limits on the PPN parameters $\\alpha_1$, $\\alpha_2$, $\\alpha_3$ and $\\xi$ will be discussed here in relation to the last two items. Tests of the strong equivalence principle (SEP) giving limits on $\\Delta$ will also be discussed due to the similar nature of the tests. These tests are null tests and it is the 90\\% confidence level limits which are quoted. In placing such limits, one wishes to know the extent to which strong field effects are contributing. Measurement of a given PPN parameter $\\hat{\\alpha}$ contains both a weak field contribution $\\alpha $ and a strong field contribution $\\alpha^{'}$ (\\cite{de92a}) \\begin{equation} \\hat{\\alpha} = \\alpha + \\alpha^{'}(c_1 + c_2 + \\cdot\\cdot\\cdot) + \\cdot\\cdot\\cdot \\end{equation} Here $c_1$, $c_2$ represent the compactness ($E_{grav} /mc^2$) of the bodies involved. For the sun, $c_i \\sim 10^{-6}$, for a neutron star $c_i \\sim 0.2$ and for a black hole $c_i \\sim 0.5$. Therefore, strong field effects are poorly constrained by solar system experiments, while pulsars provide comparable sensitivity and ease of study when compared to black holes. The cosmic microwave background (CMB) has been chosen as the absolute frame in most studies. While some recent results (\\cite{lp94}) have questioned this, it is the magnitude, not the direction of the absolute velocity ${\\bf w}$ that is most relevant; this is similar for both the CMB and Lauer \\& Postman data. One might ask whether the similar nature (i.e. upper limits from low eccentricity orbits) of the tests discussed below (which constrain $\\Delta$, $\\alpha_1$ and $\\alpha_3$) makes them degenerate. This is not the case; there are sufficient degrees of freedom and different figures of merit for each test so that different pulsars are being used for each test. ", + "conclusions": "" + }, + "9701/astro-ph9701048_arXiv.txt": { + "abstract": "We have obtained a far-ultraviolet spectrum (1150 -- 1600 \\AA) of a hot subdwarf star behind the remnant of SN~1006 with the Faint Object Spectrograph (FOS) on the {\\it Hubble Space Telescope}. The high-quality spectrum is used to test previous identifications of the strong absorption features discovered with the {\\it International Ultraviolet Explorer}. These features have FWHM $=$ 4000 ($\\pm$300) km s$^{-1}$ and are {\\it not} at the rest wavelengths of known interstellar lines, as opposed to the broader ($\\sim$8000 km s$^{-1}$ FWHM) \\ion{Fe}{2} lines from the remnant centered at zero km s$^{-1}$ in near-UV FOS spectra. We confirm that the broad absorption features are principally due to redshifted \\ion{Si}{2}, \\ion{Si}{3}, and \\ion{Si}{4} lines, which are centered at a radial velocity of 5100 ($\\pm$ 200) km s$^{-1}$. The \\ion{Si}{2} $\\lambda$1260.4 profile is asymmetric, with a nearly flat core and sharp red wing, unlike the \\ion{Si}{2} $\\lambda$1526.7 and \\ion{Si}{4} $\\lambda\\lambda$1393.8, 1402.8 profiles. One possible explanation is additional absorption from another species. Previous work has suggested that \\ion{S}{2} $\\lambda\\lambda\\lambda$ 1250.6, 1253.8, 1259.5 at a radial velocity of $\\sim$6000 km~s$^{-1}$ is responsible, but this would require a sulfur to silicon abundance ratio that is at least a factor of ten higher than expected. Another possible explanation is that the \\ion{Si}{2} and \\ion{Si}{4} profiles are intrinsically different, but this does not explain the symmetric (albeit weaker) \\ion{Si}{2} $\\lambda$1526.7 profile. ", + "introduction": "A hot subdwarf star (the ``SM star'') discovered by Schweizer \\& Middleditch (1980) lies on the far side of the remnant of SN~1006, and provides an unusual opportunity to study its ejecta in absorption. IUE observations (Wu et al. 1983; Fesen et al. 1988) revealed the presence of strong, broad \\ion{Fe}{2} absorption lines centered at zero velocity. Subsequent FOS spectra (Wu et al. 1993) were obtained at sufficient signal-to-noise and resolution to remove the narrow interstellar and stellar lines, deblend the broad \\ion{Fe}{2} lines, and determine an intrinsic \\ion{Fe}{2} profile that is suitable for comparison with model predictions for Type Ia supernovae. The \\ion{Fe}{2} velocity profile is roughly symmetric around zero km s$^{-1}$ and extends up to about $\\pm$8000 km s$^{-1}$ at the continuum. The absorption-line width and angular size of the remnant yield a lower limit to the distance of 1.9 kpc (more recent estimates based on the NW optical filament's proper motion and shock velocity yield 1.8 $\\pm$ 0.3 kpc, see Laming et al. 1996). Whereas the remnant contains approximately 0.014 $M_\\odot$ of Fe$^{+}$, the predicted mass of Fe from Type Ia supernova models (Nomoto, Thielemann, \\&\\ Yokoi 1984; H\\\"{o}flich \\& Khokhlov 1996) is about 25 times this value. However, most of the Fe in SNR 1006 could be in higher ionization states than \\ion{Fe}{2} (Hamilton \\&\\ Fesen 1988; Blair, Long, \\& Raymond 1996). The IUE observations also showed strong, broad absorption lines in the far-UV, but these features were not at the rest wavelengths of known interstellar lines. Wu et al. (1983) proposed that the outer portions of the ejecta are less chemically processed. They identified the features as \\ion{Si}{2}, \\ion{Si}{3}, and \\ion{Si}{4} lines originating from a clump in the far side of the ejecta, which is moving at a radial velocity of $\\sim$5000 km s$^{-1}$. Based on a spectrum of SN 1006 with the {\\it Hopkins Ultraviolet Telescope} (HUT), Blair et al. (1996) attribute these lines entirely to Si. However, Fesen et al. (1988) showed that the wavelength correspondences are not exact, the line strengths are not quite in the expected ratios, and some of the lines are asymmetric, and they suggested that \\ion{S}{2} at $\\sim$6000 km s$^{-1}$ and \\ion{O}{1} at $\\sim$6500 km s$^{-1}$ may also contribute to these features. Fesen et al. (1988) suggested that these features may arise from clumps of shocked material composed of intermediate mass elements in the far side of the ejecta. We have obtained high signal-to-noise spectra of the far UV lines to test these identifications, and we present the most likely interpretations of these features. A companion paper (Hamilton et al. 1996) presents a detailed physical picture of the ejecta based on one of these interpretations. ", + "conclusions": "An interesting result from the FOS observations is the large intrinsic width of the far-UV absorption lines from SNR 1006, which might be expected to resolve into individual components if the absorption arises from small knots of gas. In fact, the profiles of the deblended Si lines start at v$_{r}$ $=$ 1500 ($\\pm$500) km s$^{-1}$ and extend out to v$_{r}$ $=$ 8500 ($\\pm$500) km s$^{-1}$, which is close to the extent of the red wing of the \\ion{Fe}{2} lines (at $\\sim$8000 km s$^{-1}$). This suggests that the Si lines do not arise from a small clump or a few clumps of material, but rather in a layer of gas in bulk motion in the far side of the ejecta, and that Si and Fe may be intermixed on the far side. We note that there is substantial observational and theoretical evidence for mixing of the interior and exterior layers in a Type II supernova (SN1987A, see Bussard, Burrows, \\&\\ The 1989; Yamada \\&\\ Sato 1991). The equivalent widths of the broad absorption lines have not varied at the $\\geq$20\\% level over $\\sim$12 years of UV observation, since the first IUE observations by Wu et al. (1983). This lack of variability is also difficult to reconcile with small clumps, which have significant transverse velocities because the hot subdwarf is $\\sim$2.5$'$ from the center of the SNR. From the 12-year time interval, projected offset of the subdwarf, and observed velocities, a lower limit of 0.02 pc is obtained for the size of the absorbing region perpendicular to the line of sight. Another interesting result from the FOS spectra is that there is no evidence for strong blue-shifted counterparts to the redshifted Si lines (although there appears to be a weak depression in the $\\sim$ 1375 \\AA\\ region, just blueward of the \\ion{Si}{4} $\\lambda\\lambda$1393.8 line). The lack of blue-shifted Si lines can be explained by asymmetry in the original supernova explosion and/or asymmetry in the distribution of the interstellar medium surrounding the supernova (along with a significant interaction of the ejecta with the interstellar medium). Hamilton et al. (1996) discuss these possibilities in detail, and argue for a much lower density of the interstellar medium on the far side of the remnant compared to the rest of the remnant. The apparent excess absorptions in the 1282 \\AA\\ and 1331 \\AA\\ features are difficult to understand. They may be interpreted as \\ion{S}{2} and \\ion{O}{1} lines with centroids at higher redshifted radial velocities, but then the relative abundance of sulfur would have to be much higher than predicted from models of Type Ia supernovae. The column densities of the Si ions, determined from direct integration of the optical depths across the profiles, are N(\\ion{Si}{2}) $=$ 6.7 ($\\pm$0.6) x 10$^{14}$ cm$^{-2}$, N(\\ion{Si}{3}) $=$ 3.1 ($\\pm$0.6) x 10$^{14}$ cm$^{-2}$, and N(\\ion{Si}{4}) $=$ 3.6 ($\\pm$0.5) x 10$^{14}$ cm$^{-2}$. (The \\ion{Si}{2} value comes from \\ion{Si}{2} $\\lambda$1526.7 at 1553 \\AA, whereas Hamilton et al. [1996] obtain a higher value by assuming the entire 1282 \\AA\\ feature is due to \\ion{Si}{2} $\\lambda$1260.4.) The column densities of the \\ion{S}{2} and \\ion{O}{1} lines, estimated from the excess absorption, can be used to determine a lower limit to their abundance {\\it relative} to the total Si abundance (assuming no other ionization states for Si). The ratio of oxygen to silicon abundance is $\\geq$ 1, which is close to the expected ratio of 1 -- 2 (Nomoto et al. 1984). The ratio of sulfur to silicon abundance is $\\geq$5.5, which is a factor of ten higher than the expected ratio of $\\sim$0.5 (Nomoto et al. 1984). Thus, although the identification of the silicon lines is secure, the interpretation of excess absorption due to redshifted sulfur and oxygen lines is in doubt. Hamilton et al. (1996) discuss an alternate interpretation of the 1282 \\AA\\ feature, in which the asymmetric profile is due to shocked and unshocked components of \\ion{Si}{2} $\\lambda$1260.4 absorption. However, this does not explain the difference between the observed \\ion{Si}{2} $\\lambda$1260.4 and \\ion{Si}{2} $\\lambda$1526.7 line profiles. Since the present FOS data are insufficient to resolve this apparent discrepancy, observations at higher resolution (to fit and remove the narrow lines more accurately) and higher signal-to-noise (to get better profiles of the weak broad lines) are desirable. Improved data will also set stricter limits on blueshifted absorption features from the remnant's approaching hemisphere. Similar observations of a hot subdwarf star that is chosen to match the effective temperature and surface gravity of the SM star would also be helpful in determining an accurate continuum and removing stellar absorption features." + }, + "9701/astro-ph9701081_arXiv.txt": { + "abstract": "New high quality absorption spectra of \\jeno\\ HCO$^+$, HCN, HNC and the \\jowt\\ line of CS towards the center of the radio core of Cen A (NGC\\,5128) are presented. In addition the absorption profile of the \\jeno\\ line of H$^{13}$CO$^+$ has been detected for the first time, revealing that most absorption features have low opacities. The new HCO$^+$ spectrum allows a comparison with results obtained more than 7 years earlier. No significant change in the spectrum is found. From this remarkable result, constraints can be put on the minimum size of the radio source in the millimeter range ($>$500 AU). A comparison of abundance ratios between different absorption components suggests that the absorbing gas is essentially low density gas, with both low excitation and kinetic temperatures. Relative molecular abundances are compatible with those of the Milky Way. A parametrization of the high signal--to--noise spectrum of HCO$^+$ is presented for future reference when looking for spectral changes. ", + "introduction": "The nearby giant elliptical galaxy Cen A (NGC\\,5128) contains a warped disk of dust and dense gas. Optically the disk is seen as obscuration against a background of the old stellar population belonging to the elliptical galaxy. The presence of H$\\alpha$ emission in the disk (cf. Nicholson et al. 1992) suggest that formation of massive stars are currently taking place in the disk. The disk is also a source of molecular line emission. Several studies have shown that it contains about $3 \\times 10^{8}$\\,\\mo\\ of \\htwo. The distribution of the molecular gas has been traced by the CO emission (Eckart et al. 1990a, Quillen et al. 1992, Rydbeck et al. 1993, Wild et al. 1997). The emission extends to a galactocentric distance of approximately 1\\,kpc. The molecular gas distribution and its kinematics is consistent with a thin disk which is severly warped (Quillen et al. 1992). The main components are a ring or spiral arm at a galactocentric distance of $\\sim$800\\,pc\\ (adopting a distance of 3\\,Mpc to Cen A, which means that 1'' corresponds to 14.5\\,pc.) and a circumnuclear ring at a radius of $\\sim$100\\,pc (Israel et al. 1990, Rydbeck et al. 1993). The inner ring is seen as high velocity wings in spectra towards the center when the angular resolution is better than 25--30''. It has also been imaged with the aid of deconvolution of single dish CO(2--1) data (Rydbeck et al. 1993). The inner ring is inclined relative to the outer disk but aligned perpendicular to the inner radio jet. The molecular gas properties of the disk appears to be similar to those found in normal spiral galaxies (cf. Israel et al. 1990, Eckart et al. 1990, Wild et al. 1997). The radio core is hidden behind a large column of dense obscuring gas. The combination of a strong radio continuum source and a large column of dense gas makes the line of sight towards the center of Cen A a rich source of molecular absorption lines. The properties of the gas seen in absorption is still largely unknown. Several studies have come up with conflicting results, both concerning the location of the absorption components relative to the nucleus as well as the temperature and density of the gas. The HI absorption towards the nucleus shows three absorption components; one strong at the systemic velocity around 552\\,\\kms\\ and two redshifted ones at 596 and 609\\,\\kms, respectively. Towards the inner jet, only the main absorption at 552\\,\\kms\\ is seen (van der Hulst et al. 1983). This has been interpreted as evidence that the main 552\\,\\kms\\ line is situated far out in the disk, while the redshifted lines are situated very close to the nucleus, possibly falling in to the center. However, Seaquist \\& Bell (1990) report a detection of redshifted H$_2$CO $\\lambda$2cm absorption against the inner jet at a velocity of $\\sim$576\\,\\kms. This is not necessarily a proof against the redshifted component being situated close to the nucleus. The inner jet is seen at a projected distance of $\\sim$20'' from the core, which corresponds to $\\sim$300\\,pc. The inner molecular disk can be extended on these scales (cf. Israel et al. 1991, Rydbeck et al. 1993, Hawarden et al. 1993). The molecular absorption lines seen in the millimeter range only occurs towards the radio core of Cen A, since the inner jet has a steep radio spectrum with a completely negligible continuum flux at mm wavelengths. Molecular absorption lines seen in our Galaxy (cf. Lucas \\& Liszt 1996) and towards high redshift galaxies (Wiklind \\& Combes 1996ab, 1995, Combes \\& Wiklind 1996) almost exclusively arise in very cold gas (in terms of excitation temperature). Although the abundance ratio of HCN/HNC imply that the kinetic temperature can be in the range 10--20\\,K, the excitation temperature is comparable to the cosmic microwave background. This suggests diffuse gas with n(H$_2$) $< 10^{3}$ cm$^{-3}$. Are the molecular absorption lines seen in Cen A likewise coming from diffuse gas? Unfortunately very few multiline transitions of the same molecule have been observed. One exception is H$_2$CO, for which Seaquist \\& Bell (1990) derive an upper limit to the excitation temperature of 3.9\\,K. Several OH absorption features have been detected by van Langevelde et al. (1995), their interpretation is complicated by the presence of maser lines. Some come from diffuse gas, and some features point to dense clumps (n(H$_2$) $>$ 10$^4$ cm$^{-3}$). The three lowest rotational lines of CO have been seen in absorption, both the lines around the systemic velocity and the redshifted components. The mere detection of the CO(3--2) line (cf. Israel et al. 1991) implies that the excitation temperature is relatively high (10--20\\,K). For CO, however, the analysis is complicated by confusion with emission, especially for the two lowest transitions. In this paper we present new high quality observations of the HCO$^+$(1--0), HCN(1--0), HNC(1--0) and CS(2--1) absorption lines, as well as the previously unobserved H$^{13}$CO$^+$(1--0) transition, in order to shed some light on the physical properties and location of the molecular absorption towards the radio core of Cen A. In Sect.\\,2 we present the observations, in Sect.\\,3 we identify the different absorption components and question their possible time variations and in Sect.\\,4 we derive column densities and abundance ratios. In Sect.\\,5 we discuss the implied properties of the absorbing gas in Cen A, and its assumed distance from the center. \\begin{figure*} \\psfig{figure=fig1.ps,bbllx=10mm,bblly=60mm,bburx=205mm,bbury=205mm,width=17.5cm,angle=0} \\caption[]{Spectra of the observed transitions: HCO$^+$(1--0), H$^{13}$CO$^+$(1--0), HCN(1--0), HNC(1--0), CS(2--1) and N$_2$H$^+$(1--0). Only N$_2$H$^+$ remains undetected. The spectra have been normalized to a continuum level of unity. The velocity scale is heliocentric and the frequency definition is relativistic (see text). The velocity resolution is 0.2\\,\\kms\\ for HCO$^+$, 0.4\\,\\kms\\ for HCN, HNC and CS, and 0.6\\,\\kms\\ for H$^{13}$CO$^+$ and N$_2$H$^+$. } \\end{figure*} \\begin{table*} \\begin{flushleft} \\caption[]{Observed molecules} \\small \\begin{tabular}{l|cccc} \\hline & & & & \\\\ \\multicolumn{1}{c|}{Molecule} & \\multicolumn{1}{c}{Transition} & \\multicolumn{1}{c}{Frequency$^{a)}$} & \\multicolumn{1}{c}{$T_{\\rm cont}$} & \\multicolumn{1}{c}{Observing date} \\\\ \\multicolumn{1}{c|}{ } & \\multicolumn{1}{c}{J$\\longrightarrow$J--1} & \\multicolumn{1}{c}{GHz} & \\multicolumn{1}{c}{mK$^{b)}$} & \\multicolumn{1}{c}{ } \\\\ & & & & \\\\ \\hline & & & & \\\\ H$^{13}$CO$^+$ & 1--0 & 86.754294 & 409,355,308 & Dec95, Jul96, Aug96 \\\\ HCN & 1--0 & 88.630415 & 320 & Jul96 \\\\ HCO$^+$ & 1--0 & 89.188518 & 410,340 & Dec95, Jul96 \\\\ HNC & 1--0 & 90.663450 & 325 & Jul96 \\\\ N$_2$H$^+$ & 1--0 & 93.173500 & 314 & Dec95 \\\\ CS & 2--1 & 97.980968 & 350 & Jul96 \\\\ & & & & \\\\ \\hline \\end{tabular} \\ \\\\ a)\\ The rest--frequency of the observed molecules ($\\nu_0$ in Eqs\\,1--2) taken from Lovas (1992). \\\\ b)\\ In the $T_{\\rm A}^{*}$ temperature scale. Conversion to Jansky: 25\\,Jy/K. \\end{flushleft} \\end{table*} \\begin{figure*} \\psfig{figure=fig2.ps,bbllx=25mm,bblly=45mm,bburx=185mm,bbury=225mm,width=14.5cm,angle=0} \\caption[]{The four main lines in the Low Velocity complex is shown to the left for CS(2--1), HNC(1--0) and HCO$^+$(1--0). Also shown are the results from a 5--component gaussfit. The residuals shown above the lines, are spectral values minus the fitted values. The fift component corresponds to the High Velocity complex, shown to the right. The HV complex has been binned to a velocity resolution of 3\\,\\kms. Notice the different scales. These five components define the main absorption lines (see Table\\,2).} \\end{figure*} ", + "conclusions": "\\subsection{Type of molecular clouds} The column density of HCO$^+$ in Cen A ranges between $3.0 \\times 10^{12}$\\,\\cmsq\\ to $4.6 \\times 10^{13}$\\,\\cmsq\\ (Table\\,6). These column densities are all higher than the onset of CO self--shielding, which occurs at HCO$^+$ column densities $(1-3) \\times 10^{12}$\\,\\cmsq\\ (Lucas \\& Liszt 1996). Hence the absorption in Cen A arises in gas which has a relatively low C$^+$/C ratio. If the gas in Cen A follows the same tight correlation between $N(HCO^+)$ and $N(OH)$ as Galactic clouds (Lucas \\& Liszt 1996), our HCO$^+$ measurement imply OH column densities in the range $(0.1-2) \\times 10^{15}$\\,\\cmsq. These $N(OH)$ values are consistent with those inferred by van Langevelde et al. (1995) in Cen A. In Fig.\\,9 we plot the column density of HCO$^+$ versus the column density of HCN, HNC and CS. We also include absorption line data from Galactic diffuse clouds obtained by Lucas \\& Liszt (1993, 1994, 1996) and somewhat denser clouds seen towards Sgr\\,B2 (Greaves \\& Nyman 1996). The latter clouds are in some cases blended with gas close to the Galactic center. We have multiplied the HCO$^+$ column densities of Lucas \\& Liszt with a factor 1.4 in order to correct for the use of different values of the electric dipole moment. The dotted line in the figure is not a fit to the data but represents a one--to--one correspondance between the column densities. The gas in Cen A follows the same correlation as that of Galactic molecular gas. We have used an excitation temperature of 5\\,K when deriving column densities, while Lucas \\& Liszt and Greaves \\& Nyman have used 2.76\\,K. Since the molecules have almost exactly the same dependence on $T_{\\rm x}$, this has no influence on the result. Although the excitation temperature is likely to be low in the molecular gas seen in absorption, the HCN/HNC ratios implies that the kinetic temperature is rather high. The formation of HCN and HNC depends on the gas temperature (cf. Irvine et al. 1987), with HCN preferentially being formed in warm gas on behalf of its isotopomer HNC. In the LV complex we find HCN/HNC ratios ranging between 3--7, while the HV complex has a ratio of 2.2. This gives a kinetic temperature of the LV complex of 20--30\\,K, while the HV gas is characterized by a kinetic temperature of $\\la$10\\,K. A more detailed comparison of the 5 main absorption components as defined in Table\\,2 and 6, shows that component 1 and 3 (in the LV complex) shows similar behaviour in their abundance ratios (see Table\\,6). The differences between the components are, however, relatively small and do not show up as significant deviations in Fig.\\,9. The ratios of the hyperfine components of HCN (Table\\,4) indicate that component 1 and 3 are close to the LTE values; R$_{12}=0.6$, R$_{02}=0.2$, while component 2 and 4 deviates from LTE. Here we have to keep in mind that in the decomposition of the hyperfine lines some components suffer from near overlap and that the F$=$0--1 line of component 4 is an upper limit. Nevertheless, it is tantalizing that the two components with close to LTE ratios also show similar abundance ratios, while the other two components (plus the HV complex) differs by factors of 2--3 from each other. Anomalous excitation of the HCN hyperfine lines is often seen in Galactic molecular clouds (cf. Guilloteau \\& Baudry 1981, Walmsley et al. 1982), with the R$_{12}$ ratio lower than LTE values in warm clouds and the R$_{01}$ ratio lower than LTE values in cold clouds. In dense regions, where the HCN(1--0) line thermalizes, both ratios tend to unity. This latter case appears to be the case for component no. 2 in Cen A, at least for the R$_{12}$ ratio. In Fig.\\,10 we compare the HCO$^+$(1--0) spectrum with the HI absorption obtained with the VLA (van der Hulst et al. 1983). The HCO$^+$ spectrum has been binned to the same velocity resolution as the HI data, 6.2\\,\\kms. A 3--component gaussian fit to the HCO$^+$ spectrum and the residual is also shown. The appearance of the spectra agree quite well, even though the redshifted HCO$^+$ is more spread out in velocity than the HI. Column densities and N(HI)/N(HCO$^+$) ratios are given in Table\\,8. The component at $\\sim$580\\,\\kms\\ appears to have a higher molecular gas fraction than the other two by a factor of $\\sim$3. In summary, all of the molecular gas components seen in absorption towards the nucleus of Cen A have a chemistry similar to that of Galactic diffuse molecular gas. Only the column densities are higher in Cen A than in typical Galactic diffuse clouds, which could be due to unresolved line components (i.e. more clouds) in the line of sight through the edge--on disk. The narrow HCO$^+$ lines seen in the HV complex reveal a relatively diffuse gas with lower abundances and a low kinetic temperature (T$_{\\rm k} \\le$ 10K). \\begin{figure} \\psfig{figure=fig10.ps,width=8.5cm,angle=-90} \\caption[]{The HCO$^+$(1--0) spectrum binned to a velocity resolution of 6.2\\,\\kms, which is the same as the 21cm HI spectrum obtained by van der Hulst et al. (1983). Also shown is a 3--component gaussian fit and the residual spectrum.} \\end{figure} \\begin{table} \\begin{flushleft} \\caption[]{Column densities for HCO$^{+}$ and HI$^{a)}$.} \\scriptsize \\begin{tabular}{ccccc} \\hline & & & & \\\\ \\multicolumn{2}{c}{Comp.$^{b)}$} & \\multicolumn{1}{c}{N(HCO$^+$)} & \\multicolumn{1}{c}{N(HI)} & \\multicolumn{1}{c}{N(HCO$^+$)/N(HI)} \\\\ & & & & \\\\ \\multicolumn{1}{c}{V(HCO$^+$)} & \\multicolumn{1}{c}{V(HI)} & \\multicolumn{1}{c}{cm$^{-2}$} & \\multicolumn{1}{c}{cm$^{-2}$} & \\multicolumn{1}{c}{$10^{-8}$} \\\\ & & & & \\\\ \\hline & & & & \\\\ 550 & 553 & 13.683 & 21.415 & 0.5 \\\\ 582 & 576 & 13.098 & 21.279 & 1.5 \\\\ 608 & 596 & 13.063 & 20.708 & 0.4 \\\\ & & & & \\\\ \\hline \\end{tabular} \\ \\\\ Column densites are given as $\\log(N/{\\rm cm}^{-2})$.\\\\ a)\\ HCO$^+$ is smoothed to a velocity resolution of 6.2\\,\\kms. \\\\ Column densities for HI derived assuming $T_{\\rm s}=100$\\,K. \\\\ b)\\ Center velocities for the HCO$^+$ and HI lines (see Fig.\\,10). \\end{flushleft} \\end{table} \\subsection{Location of the absorbing molecular gas} The presence of two absorption complexes in Cen A, one at the systemic velocity and one redshifted relative to the systemic velocity, has led to speculations that the redshifted component arises in gas falling into the nucleus of Cen A, possibly feeding a supermassive black hole (cf. van der Hulst et al. 1983, van Gorkom et al. 1989). One argument is that the redshifted HI component is only seen towards the radio core and not against the inner jet (van der Hulst et al. 1983). The detection of $\\lambda$2cm H$_2$CO in absorption against the inner jet (Seaquist \\& Bell 1990) is not necessarily a proof against the infall hypothesis, if the inner circumnuclear molecular disk is extended on scales of $\\sim$300\\,pc. On the other hand, the lack of redshifted HI absorption against the inner jet could also mean that while the size of the HI component at the systemic velocity is $\\ga$300\\,pc (in order to cover both the jet and the core), the extent of the redshifted gas is smaller. HI absorption has been detected in 9 elliptical galaxies (van Gorkom et al. 1989, Mirabel 1990 and references therein). In all cases the absorption is redshifted with respect to the systemic velocity. It is not clear whether this preponderance of redshifted absorption reflects a true infall of gas or if it results from a systematic offset in the systemic velocities towards the blue. Such systematic errors are known to exist due to outflow of emission line gas, which is often used to derive the systemic velocity. In spiral galaxies the situation is different. Here HI absorption is often seen as a single broad line, extending into both the blue-- and redshifted sides around the systemic velocity (cf. Dickey 1986). With higher spatial resolution, the broad HI component is decomposed into a narrow one, shifting across the finite extent of the background radio source, over all velocities defined by the rotation of the disk (Koribalski et al. 1993). These absorptions originate in fast rotating circumnuclear disks or rings, of size $\\sim$200\\,pc, and with velocities $\\sim$200\\,\\kms\\ (e.g. NGC 253, 660, 1808, 3079, 4945, Milky Way). \\medskip The same phenomenon could be occuring in the center of Cen A, where the presence of a nuclear disk or ring of $\\sim$100\\,pc radius and rotating with a velocity of $\\sim$220\\,\\kms\\ has been established (Rydbeck et al. 1993). High rotational velocities at small galactocentric distances are naturally occuring in elliptical galaxies due to the high mass concentration towards the center (cf. Hernquist 1990) and non--circular orbits for the gas can be generated through non-axisymmetric gravitational instabilities (e.g. bar) or in the form of tri--axiality of the elliptical galaxy itself. Since the angular extent of the continuum source in Cen A is $\\la$2\\,mas (6000\\,AU), either blue or redshifted gas, depending on the orientation of the orbits, can be seen in absorption. This situation is reminiscent of what happens in the center of the Milky Way. Here the velocities are more strongly non--circular, maybe because the central bar is oriented at 20--30$^\\circ$ from the Sun line of sight (e.g. Blitz \\& Spergel 1991, Weinberg 1992). HCO$^+$ absorption in front of the central continuum source SgrA reveals a broad component between $-$210 to $-$110\\,kms, interpreted as coming from the $\\sim$200\\,pc nuclear disk, and four narrower features (at $-$51, $-$30, $-$2 and $+$32\\,\\kms) corresponding to known spiral arms in the galactic disk (Linke et al. 1981). The broad component at negative velocities is also seen in absorption in front of SgrB2, in HCO$^+$ as well as H$^{13}$CO$^+$ and HCN (Linke et al. 1981, Greaves \\& Nyman 1996). The size of the millimeter continuum source SgrA has recently been determined through VLBI at 3 and 7 mm (Krichbaum et al. 1994), and is 0.33 and 0.75 mas respectively. At those frequencies, the interstellar scattering becomes negligible, so these figures are believed to be the actual source sizes (Krichbaum et al. 1994). Since these source sizes correspond to $\\sim$5 AU at the Galactic Center distance, the HCO$^+$ absorption features prove that the apparent line of sight velocity dispersion can be quite high in a typical edge--on nuclear disk, due to the accumulation on the line of sight of differential non--circular motions. The emitting molecular gas in Cen A is confined to the center region, consisting of the nuclear disk or ring at a radius of $\\sim$100\\,pc and an outer ring (or spiral arm) at $\\sim$750\\,pc. The absorption complexes are likely to be associated with these features, but since absorption is sensitive to diffuse and low excitation molecular gas which is generally not seen in emission (e.g. Lucas \\& Liszt 1996), it is possible that some intervening molecular gas detected in absorption is associated with the larger HI disk extending to 7\\,kpc (Schiminovich et al 1994). As we have seen in the previous analysis, the molecular gas in both LV and HV components is diffuse, with abundance ratios similar to Galactic values. The main difference is that gas in the LV components has a higher kinetic temperature than the HV components. We cannot differentiate between inner and outer molecular gas from this alone. There are, however, two facts which suggest that the HV components are associated with the nuclear gas and the LV components with the outer disk: (1) the LV components are close to the systemic velocity and, (2) whereas the LV components extend between 540--556\\,\\kms, the HV components are spread out between 576--640\\,\\kms, $\\sim$4 times larger. Although the rotational velocities are similar for the inner and outer molecular gas (Rydbeck et al. 1993), the inner region has a larger velocity gradient and this gas seen in absorption should be spread over a larger velocity interval. \\medskip The blue-shifted `wing', from 500 to 540\\,\\kms, has not been considered above. This feature is seen at a very low level, and depends on the subtracted emission profile shape, and should be viewed with caution. Could this gas comes from inside the cavity delineated by the circumnuclear ring? There is a constraint on the radius where molecules can subsist, around a luminous ionizing X--ray source. Maloney et al. (1994) derived an effective ionization parameter $\\xi_{eff}$: $$ \\xi_{eff} = 1.1 \\times 10^{-2} L_{44}/(n_9 r_{pc}^2 N_{22}^{0.9}) $$ where $L_{44}$= L$_x/10^{44}$erg s$^{-1}$, $n_9 =$ n(H$_2$)/10$^9$ cm$^{-3}$ and $r_{\\rm pc}$ is the distance from the X--ray source in parsecs. The gas will not become substantially molecular unless the effective ionization parameter is smaller than 10$^{-3}$. From ROSAT HRI measurements, D\\\"obereiner et al (1996) have estimated an X-ray luminosity (0.1-2.4 kev) for the Cen A nucleus of L$_{\\rm x}= 3 \\times 10^{41}$ erg s$^{-1}$ in 1994, after correcting for absorption. For diffuse gas with n(H$_2$)$\\sim 3 \\times 10^{3}$\\,\\cmcb, and a column density N(H$_2$) $\\la$ 10$^{22}$ cm$^{-2}$, the minimum distance from the center is $\\sim$100\\,pc. For molecular gas to exist at smaller galactocentric distances in Cen A, both the volume and column densities must be considerably higher. It is therefore likely that the molecular gas corresponding to the blue--shifted wing is at least at the distance of the circumnuclear ring. It is possible that this wing corresponds to the 750\\,pc component, which should also possess non--circular motions. The orientation of orbits in a tumbling non--axisymmetric component change by 90$^\\circ$ at each resonance (e.g. Contopoulos \\& Grosbol 1989). Along the nucleus line of sight, it is therefore possible that this blue--shift component corresponds to elliptical streamlines perpendicular to that in the circumnuclear ring." + }, + "9701/astro-ph9701191_arXiv.txt": { + "abstract": "{\\baselineskip 0.4cm The {\\it COBE} satellite has provided the only comprehensive multi-frequency full-sky observations of the microwave sky available today. Assessment of the observations requires a detailed likelihood analysis to extract the maximum amount of information present in the noisy data. I present a specific method for estimating the CMB anisotropy power spectrum independent of any assumptions about the underlying cosmology, and then use standard image processing techniques to generate the most revealing corresponding maps of the signal. The consistency of the data at the three available frequencies provides strong support to the assertion that we are being provided with our first glimpse of the last scattering surface. } ", + "introduction": "Operations of the Differential Microwave Radiometer instruments (DMR), the last active experiment on board NASA's COsmic Background Explorer ({\\it COBE}) satellite, were terminated in December 1993, concluding four years of measurements of anisotropy of the cosmic microwave background (CMB) radiation. The final product of the DMR-team work --- full sky maps of the microwave sky at 31.5, 53, and 90 GHz --- were released to the astronomical community in January 1996. A generic description of the final DMR data set and brief summary of the results of the DMR-team analysis were given in \\cite{b96}. The issues of prime concern in the DMR-team work included the following: \\noindent 1) modeling and removal of identified systematic artifacts from the data (see \\cite{ksys}) to produce sky maps suitable for cosmological studies, \\noindent 2) studies of potential signal contamination by the Galaxy \\cite{kgal}, and/or the extragalactic sources \\cite{extrag}, \\noindent 3) testing of the hypothesis of primordial origin of the measured CMB anisotropy \\cite{rms} by evaluation of the frequency dependence of the $(\\delta T/T)_{rms}$ in the sky maps, \\noindent 4) testing of the hypothesis of gaussianity of statistics of the measured anisotropies \\cite{kgauss}, \\noindent 5) evaluation and analysis of the auto- and cross-correlation functions of the sky maps \\cite{hcorr}, \\noindent 6) determination of the angular power spectrum of the CMB anisotropy \\cite{g96}, \\cite{hsp}, \\cite{w96}. \\noindent Other analyses of the 4 year data presented so far include \\cite{bw}, \\cite{teg1}, and Bond and Jaffe in these Proceedings. \\begin{figure}[t] \\psfig{file=dwidma_315390_20g_new.ps,width=6.7in,angle=90} \\caption{Power spectra of the DMR (galactic frame, extended Galaxy cut, 3881 pixels) difference maps, (A-B)/2, at each frequency of observations, for each yearly sky maps (top row, stars --- year one, squares - year 2, triangles --- year 3, diamonds --- year 4), for two 2-year sky maps (middle row, stars --- years 1 and 2, diamonds --- years 3 and 4), and for the final four-year product (bottom row, diamonds). Bottom row plots include the 68, 95, and 99\\% confidence regions (heavy to light grey) from Monte Carlo simulations of instrumental noise in the sky maps.} \\end{figure} In this contribution I will focus on two aspects of the interpretation and visualisation of the \\cobedmr % data: 1) derivation of the angular power spectrum of the CMB anisotropy using a data reduction method which is {\\it independent} of the cosmological model, and 2) linear filtering of the sky maps. Both these goals are pursued within the mathematical framework for analysis of the DMR data which was originally described in \\cite{g94a}, and \\cite{g94b}. I will briefly describe the DMR data, address the necessary technical points, and proceed to presentation of the results and discussion. \\begin{figure}[t] \\psfig{file=widma_315390_20g_new.ps,width=6.7in,angle=90} \\caption{Power spectra of the DMR (galactic frame, extended Galaxy cut, 3881 pixels) inverse-noise-variance weighted sum maps, at each frequency of observations, for all yearly, 2-year, and 4-year sky maps (same coding as in Figure 1.) Variance weighting of the A and B sides makes the yearly and bi-yearly effective noise levels in the 31GHz map more consistent with one another than in the corresponding difference maps (with fifty-fifty noise contributions from the A and B sides), and the effective noise level in the sum maps somewhat lower than in the corresponding difference maps at all three frequencies. } \\end{figure} ", + "conclusions": "" + }, + "9701/astro-ph9701228_arXiv.txt": { + "abstract": " ", + "introduction": "Large scale structure (LSS) study of Big Bang cosmology tries to explain how an initially flat or smooth 3-dimensional surface described by the Robertson-Walker metric evolved into a wrinkled one. In terms of density and velocity fields, it explains how an initially homogeneous and Hubble-expanding mass distribution evolved into its present inhomogeneous state. It is generally believed that LSS was initiated by fluctuations formed at the early universe, and that the subsequent clustering was brought about by gravitational interaction between baryonic and dark matter (Kolb \\& Turner 1989). As a result, like the physics of dynamical critical phenomena, turbulence, and multiparticle production in high energy collisions, problems in LSS are typical of structure formation due to stochastic forces and non-linear coupling (Berera \\& Fang 1994, Barbero et al 1996). The cosmic mass (or number) density distribution ${\\rho(x)}$ can be mathematically treated as a homogeneous random field. Traditionally, the statistics, kinetics and dynamics in LSS are represented by the Fourier expansion of the density field, $|\\rho(k)|$. For instance, the behavior of the LSS in scale space can effectively be described by the power spectrum of perturbations $P(k) = |\\rho(k)|^2$. In the case where the homogeneous random field $\\rho(x)$ is Gaussian, all statistical features of $\\rho(x)$ can be completely determined by the amplitude of the Fourier coefficients. In other words, the two-point correlation function, or its Fourier counterpart the power spectrum, are enough to describe the formation and evolution of the LSS. However, the dynamics of LSS, such as clustering given by gravitational instability, is non-linear. Even if the field $\\rho(x)$ is initially Gaussian, the evolved density field will be highly non-Gaussian. To describe the dynamics of the LSS knowledge of the phase of the Fourier coefficients $\\rho(k)$ is essential. As is well known, it is difficult, even practically impossible, to find information about phase of the Fourier coefficients as soon as there is some computational noise (Farge 1992). This lack of information makes the description of LSS incomplete. Even in the case where the phases are detectable, the pictures in physical space, ${\\rho(x)}$, and the Fourier space, $\\rho(k)$ are separated. From the former we can only see the scales of the structures, but not the positions of the considered structures, and {\\it vice versa} from the later. It has been felt for some time that the separate descriptions between Fourier (scale) and physical (position) spaces may lead to missing key physics. In order to resolve this problem, methods of space-scale-decomposition (SSD) which might provide information about the phase (or position) and scale of the considered structures have been developed. The possibility of simultaneously localizing in both frequency (scale) and time (position) is not new in physics. Anybody who listens to music knows that they, at any time, can resolve the frequency spectrum. The problem of how to perform this time resolution is also not new in physics. Wigner functions in quantum mechanics, and the Gabor transform (Fourier transform on finite domain) were early approaches. Speaking simply, SSD represents a density field as a superposition of density perturbations localized in both physical and scale spaces. For instance, identification of clusters from a galaxy distribution by {\\em eyes} is a SSD. Generally, all methods of identifying clusters and groups from surveys of galaxies or samples of N-body simulation are SSD. One can list several popular SSDs in cosmology as follows: smoothing by a window function, or filtering technique; percolation; the friend-to-friend algorithm; count in cells (CIC). A common problem of most of the above mentioned SSDs is that the bases, or representations, given by these methods is incomplete. Unlike the Fourier representation $\\rho(k)$, these SSDs lose information contained in the density field $\\rho(r)$. For instance, one can completely reconstruct the density field $\\rho(x)$ by the Fourier coefficients $\\rho(k)$, but cannot do the same using window filters, CIC, percolation etc. All these SSDs are, directly or indirectly, the precursors to the DWT (Discrete Wavelet Transform). The DWT is also a SSD, but is based on bases sets which are orthogonal and complete. The DWT is invertible and admissible making possible a complete representation of LSS without losing information. Unlike the Fourier bases (the trigonometric functions) which are inherently nonlocal, the DWT bases have limited spatial support. The DWT allows for an orthogonal and complete projection on modes localized in both physical and space spaces and makes possible a multiscale resolution. Moreover, the orthogonal bases of the DWT are obtained by (space) translation and (scale) dilation of one scale function (Meyer 1992, 1993; Daubechies 1992). They are self-similar. This translation-dilation procedure allows for an optimal compromise: the wavelet transform gives very good spatial resolution on small scales, and very good scale resolution on large scales. Therefore, the DWT is able to resolve an arbitrary density field simultaneously in terms of its position variable and its conjugate counterpart in Fourier space (wavenumber or scale) up to the limit of uncertainty principle. There have been attempts to use the continuous wavelet transform (CWT) to analyze LSS (Slezak, Bijaoui \\& Mars 1990; Escalera \\& Mazure 1992; Escalera, Slezak \\& Mazure 1992; Martinez, Paredes \\& Saar 1993). However, since 1992 it has become clear that the CWT is a {\\it poor} or even {\\it impossible} method to use as a reasonable SSD. The difference between CWT and DWT is mathematically essential, unlike the case for the Fourier transform, for which the continuous-discrete difference is only technical (Yamada \\& Ohkitani 1991; Farge 1992; Greiner, Lipa \\& Carruthers 1995). These properties of the DWT make it unique among the various SSD methods. One can expect that some statistical and dynamical features of LSS can easily, and in fact {\\it only}, be described by the DWT representation. The DWT study of LSS now is in a very preliminary stage. Nevertheless, results have shown that the DWT can reveal aspects of LSS behavior which have not been seen by traditional methods (Pando \\& Fang 1995, 1996a, 1996b; Huang et al 1996). These DWT-represented features have also been found to be effective for discriminating among models of LSS formation. As we will show the DWT opens a new dimension in the study of the statistics and dynamics of the LSS. \\setcounter{enumi}{2} \\setcounter{equation}{0} ", + "conclusions": "" + }, + "9701/astro-ph9701122_arXiv.txt": { + "abstract": "The results of the 0.1-10 keV spectral analysis of IC 5063 (Narrow Line Radio Galaxy) and NGC 3998 (LINER) are presented. In both cases there is evidence of a FeK$\\alpha$ emission line, but its width is poorly constrained. The highly absorbed spectrum of IC 5063 is well described by a double power law model. The X-ray luminosity of a few $\\times$ 10$^{43}$ ergs s$^{-1}$ reveals a Sey 1 nucleus seen nearly edge-on, as confirmed by optical and IR studies. NGC 3998 is well fitted by a power law model ($\\Gamma$$\\sim$1.9) without reflection and weak absorption. The weak X-ray emission detected by ASCA and ROSAT suggests a low-luminosity AGN or an active nucleus in a low-accretion period. ", + "introduction": "IC 5063 is classified as a Narrow Line Radio Galaxy (z=0.011) (Caldweel \\& Phillips 1981). IC 5063 was observed with the gas imaging spectrometer (GIS) and solid state spectrometer (SIS) on the ASCA satellite (Tanaka et al. 1994) in April 1994 and with the ROSAT PSPC (Pfeffermann et al. 1987) from November 1991 to April 1992. After applaying standard criteria to the data, about 19 Ks of useful data for ROSAT and $\\sim$ 21(24) Ks for each SIS(GIS) detector were collected.\\\\ The combined analysis of ASCA and ROSAT data reveals a soft X-ray excess, a hard continuum and a reflection component. A double power law model plus a reflection component (fixed to cover 2$\\pi$ steradiants at the source) provides a good fit to the ROSAT and ASCA data, but the spectral slope of the soft power law ($\\Gamma_{\\rm S}$=2.16$^{+0.29}_{-0.31}$, errors at 90\\%, $\\Delta$$\\chi^{2}$=2.71), steeper than the hard ($\\Gamma_{\\rm H}$=1.70$^{+0.18}_{-0.19}$), indicates a more complex soft X-ray emission. A $L_{\\rm soft}$/$L_{\\rm hard}$ $\\sim$1\\% ratio is obtained by a partial covering model. A bremsstrahlung model for the soft emission is equally good (see table 1). The broad H$\\alpha$ line in polarized flux (FWZI$\\sim$6000-9000 Km s$^{-1}$, Axon et al. 1994) and the anisotropy of the radiation field (``X'' geometry), with an opening angle of $\\sim$50$^{\\circ}$ (Colina et al. 1991), suggest that the observed soft X-ray component may be emission scattered along the line of sight from ionized gas above an absorbing torus. The column density value, $N_{\\rm H}$=2.15$\\pm{0.18}$ $\\times$ 10$^{23}$ cm$^{-2}$, and the dust lane across the galaxy, likely responsible for the IR emission (Axon et al. 1982), agree with a nearly edge-on configuration. The 2-10 keV luminosity L$\\sim$2$\\times$10$^{43}$ ergs s$^{-1}$ indicates a Sey1-like active nucleus for IC 5063. Strong evidence for the FeK$\\alpha$ line is given only by the SIS1 detector. The line parameters are given in table 2.\\\\ \\centerline{\\bf Tab. 1 - Spectral parameters} \\begin{table}[h] \\hspace{0.80cm} % \\begin{tabular}{|l|c|c|c|c|c|} \\hline \\hline source&model&$\\Gamma_{\\rm S}$/kT&$\\Gamma_{\\rm H}$&$N_{\\rm H}$&$\\chi^{2}$/dof\\\\ \\hline {\\bf IC5063}&po+po&2.16$^{+0.29}_{-0.31}$&1.70$^{+0.18}_{-0.19}$&2.15$\\pm{0.18}$ $\\times$10$^{23}$&260/263\\\\ \\dotfill&brem+po&1.39$^{+1.02}_{-0.42}$&1.66$^{+0.20}_{-0.18}$& 2.07$^{+0.20}_{-0.19}$ $\\times$10$^{23}$&262/263\\\\ \\hline {\\bf NGC3998}&po&\\dotfill&1.88$\\pm{0.02}$&$\\sim$$N_{{\\rm H}_{\\rm gal}}$ &1036/688\\\\ \\hline \\end{tabular} \\end{table} ", + "conclusions": "" + }, + "9701/astro-ph9701070_arXiv.txt": { + "abstract": "Analysing the weak lensing distortions of the images of faint background galaxies provides a means to constrain the mass distribution of cluster galaxies and potentially to test the extent of their dark matter halos as a function of the density of the environment. Here I describe simulations of observational data and present a maximum likelihood method to infer the average properties of an ensemble of cluster galaxies. ", + "introduction": "Measurements of the rotation curves of spiral galaxies indicate that they are embedded in massive dark matter halos. The deflection of light rays through the gravitational action of mass concentrations, usually called gravitational lensing, provides a way to obtain information about the mass distribution of galaxies at radial distances from their centre where there are no more luminous test particles to probe the gravitational potential. The light deflection causes small distortions of the images of faint background galaxies. Recent statistical analyses (Brainerd et al. 1996, Griffiths et al. 1996) of these weak distortion effects suggest that the dark galaxy halos are indeed rather extended, as some popular theories of structure formation predict them to be. During the formation of galaxy clusters the extended halos of galaxies may be stripped off due to tidal forces of the cluster potential or during encounters with other galaxies. Ultimately the individual galaxy halos should merge and form a global cluster halo. In this contribution I discuss how this merging picture could be tested observationally by exploiting the weak lensing effects. The distortions of the images of background galaxies produced by massive galaxy clusters are strong enough to allow a parameter-free reconstruction of the clusters' surface mass density, and several algorithms have been developed for this purpose (e.g. Kaiser and Squires 1993, Seitz and Schneider 1995, 1996). The smoothing length which has to be implemented in these techniques, however, is larger than galaxy scales, i.e., the amount of information available does not suffice to reconstruct cluster galaxies individually. Therefore, one has to superpose the effects of a large number of galaxies statistically in order to infer the average properties of an ensemble of galaxies. Section\\,\\ref{simulations} presents simulations of a galaxy cluster which are sufficiently realistic for the purposes of this work, and demonstrates how individual galaxies modify the distortion pattern of a smooth cluster mass distribution. Section\\,\\ref{method} discusses a maximum likelihood method for constraining the mass distribution of cluster galaxies, and Sect.\\,\\ref{results} presents results of the simulations. Finally, in Sect.\\,\\ref{prospects} some suggestions for refining the simulations are mentioned, and observational prospects are discussed. A closely related work was recently published by Natarajan and Kneib (1997); in contrast to their maximum likelihood method, the mass profile of the cluster is not assumed to be known but is reconstructed from image distortions as mentioned above. ", + "conclusions": "" + }, + "9701/astro-ph9701185_arXiv.txt": { + "abstract": "We have obtained deep, near-IR images of nearly 100 host galaxies of nearby quasars and Seyferts. We find the near-IR light to be a good tracer of luminous mass in these galaxies. For the most luminous quasars there is a correlation between the maximum allowed B-band nuclear luminosity and the host galaxy mass, a ``luminosity/host-mass limit''. Comparing our images with images from HST, we find that the hosts of these very luminous quasars are likely early type galaxies, even for radio-quiet objects whose lower-luminosity counterparts traditionally live in spirals. We speculate that the luminosity/host-mass limit represents a physical limit on the size of black hole that can exist in a given galaxy spheroid mass. We discuss the promises of NICMOS for detecting the hosts of luminous quasars. ", + "introduction": "Although most of the attention given to AGN is concentrated on the ``N,'' there are compelling reasons to understand the ``G.'' The central engine and the host galaxy must influence each other, and the exact connections hold crucial clues for understanding the quasar phenomenon. Moreover, it is plausible that nuclear activity has played a role in the evolution of a significant fraction of all galaxies; Seyferts account for $\\gtrsim10\\%$ of galaxies today (Maiolino \\& Rieke 1995; Ho 1996), and AGN were even more important in the past. Therefore, to understand the evolution of galaxies, we must understand the host galaxies of AGN. By now it is well established that the redshift range $2\\lesssim z \\lesssim3$ represents a critical period in the evolution of both ``normal'' galaxies and quasars. It is likely that galaxies at that epoch were starting to turn their gas into stellar disks. The mass in neutral hydrogen gas in damped Lyman$-\\alpha$ absorbers at that redshift is comparable to the mass in disk stars today and shows strong evolution since that time (Lanzetta \\etal~1995; Storrie-Lombardi \\etal~1996). Furthermore, a photometric-redshift analysis of the Hubble Deep Field shows that the luminosity density from star-forming galaxies peaks near that same redshift (Sawicki \\etal~1996). The AGN luminosity function shows a similarly strong evolution; if described in terms of luminosity evolution (but see Wisotzki et al. in these proceedings), the ``characteristic'' luminosity of quasars increases as $\\sim(1+z)^{3.4}$ to $z\\sim2$ and flattens between $2 30\\deg$, surveying about 786 square degrees of sky with the Automated Plate Measuring (APM) system at Cambridge.\\footnote{The APM is a National Astronomy Facility, at the Institute of Astronomy, operated by the Royal Greenwich Observatory. A general description of the APM facility is given by \\cite{kibblewhite84}.} We have identified 693 galaxies, most previously uncataloged and most with central surface brightness $\\mu_B(0) > 22~\\magsq$. The complete catalog of this survey appears in \\cite{impey96a} (Paper I). The selection effects and completeness corrections for the survey are analyzed in detail in \\cite{sprayberry96a} (Paper~II). In this paper, we present the luminosity function for LSB galaxies from the APM survey and compare that luminosity function to those obtained from the CfA redshift survey. We also review suggestions by \\cite{phillipps90}, \\cite{mcgaugh94b}, \\cite{mcleod94}, and \\cite{ferguson94a} that LSB galaxies might account at least partially for the large numbers of faint blue galaxies seen in deep surveys. Section~\\ref{sec:smp} describes the survey data and presents the samples used for determining the luminosity function and the corrections applied to those samples. Section~\\ref{sec:mth} covers the methods used to develop the luminosity functions. Section~\\ref{sec:res} presents the luminosity functions and compares the results to those obtained from the CfA redshift survey. Section~\\ref{sec:imp} reviews the consequences of this LSB luminosity function for the general field luminosity function and for the question of local counterparts to the faint blue galaxies. Finally, Section~\\ref{sec:cnc} summarizes our conclusions. Throughout this paper, we assume \\hubble. Also, all magnitudes and surface brightnesses used here are in the Johnson $B$ band. ", + "conclusions": "\\label{sec:cnc} We have estimated a luminosity function for galaxies with surface brightnesses fainter than $\\mu(0) = 22.0 \\, \\magsq$, which is the approximate faint limit of $\\mu(0)$ for galaxies covered by the CfA Redshift Survey. We find that this LSB LF has a steeply rising tail at low luminosities($\\alpha = 1.42$), comparable to that found by \\cite{marzke94a} for galaxy types $8 \\leq T \\leq 10$. The LSB LF has a normalization lower than that found for the overall CfA survey, but much higher than that found for types $8 \\leq T \\leq 10$. Thus estimates of the total population of local galaxies based on the CfA survey are missing at least one-third of the total number of galaxies due to surface brightness selection bias. These previously unaccounted-for LSB galaxies can help considerably to resolve the apparent difference between estimates of the local population and the large numbers of faint blue galaxies observed at moderate redshift." + }, + "9701/astro-ph9701179_arXiv.txt": { + "abstract": "We present and discuss the X-ray luminosity function (XLF) of the ROSAT Brightest Cluster sample (BCS), an X-ray flux limited sample of clusters of galaxies in the northern hemisphere compiled from ROSAT All-Sky Survey data. The BCS allows the local cluster XLF to be determined with unprecedented accuracy over almost three decades in X-ray luminosity and provides an important reference for searches for cluster evolution at higher redshifts. We find the significance of evolution in both the XLF amplitude and in the characteristic cluster luminosity $L_X^{\\star}$ to be less than $1.8\\sigma$ within the redshift range covered by our sample thereby disproving previous claims of strong evolution within $z\\la 0.2$. ", + "introduction": "The ROSAT Brightest Cluster sample (BCS, Ebeling et al.\\markcite{bcsi} 1996b, hereafter Paper I) is a 90\\% complete, flux limited sample of the 199 X-ray brightest clusters of galaxies in the northern hemisphere ($\\delta \\geq 0^{\\circ}$), at high Galactic latitudes ($|b| \\geq 20^{\\circ}$, with redshifts $z \\leq 0.3$, fluxes higher than $4.45\\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$ and luminosities higher than $5\\times 10^{42}$ erg s$^{-1}$ in the 0.1--2.4 keV band. Second in size only to the XBACs sample of Ebeling et al.\\markcite{xbacs} (1996a), the BCS is one of the largest statistical cluster samples compiled at X-ray wavelengths to date. It is the only large-scale sample available today that is not only X-ray {\\em flux limited}\\/ but also X-ray {\\em selected} in the sense that the BCS, unlike the XBACs, is not limited to systems initially found in optical surveys but contains clusters selected by their X-ray properties only. The BCS thus represents an ideal sample for studies of the formation, distribution and evolution of structure on the largest metric and mass scales. Providing important constraints on the cosmological parameters governing cluster evolution, the X-ray luminosity function (XLF) of clusters of galaxies represents a particularly vital statistic in this context. Several previous studies based on much smaller samples of typically 50 clusters found the evolution in the cluster X-ray luminosity function to be `negative' in the sense that X-ray luminous clusters are more numerous now than they were in the past\\markcite{edge,gioia,henry,david} (Edge et al.\\ 1990; Gioia et al.\\ 1990; Henry et al.\\ 1992; David et al.\\ 1993). They were, however, not only in conflict with other studies, which found no evidence for cluster evolution (e.g.,\\markcite{kowalski} Kowalski et al.\\ 1984), but also somewhat inconsistent among themselves. The strong evolution seen by Edge et al.\\ (1990) in their sample of 46 X-ray bright clusters at high Galactic latitude and $z \\leq 0.18$ is not present in the first two redshift bins (44 clusters at $0.14\\leq z \\leq 0.3$) of the sample of Gioia et al.\\markcite{gioia} (1990) and Henry et al.\\markcite{henry} (1992) who find significant evolution only at $z>0.3$. More recently, two studies found no sign of evolution at all in the XLF of samples of Abell and ACO clusters at $z\\le 0.36$ \\markcite{briel} (Briel \\& Henry 1993) and $z\\le 0.15$ \\markcite{burg} (Burg et al.\\ 1994), respectively. However, these samples were neither X-ray selected nor X-ray flux limited and may thus not be fair representations of the cluster population in general. At considerably higher redshifts ($z>0.5$) on the other hand, studies based on yet smaller samples observed in deep X-ray pointings \\markcite{bower,castander1,castander2} (Bower et al.\\ 1994; Castander et al.\\ 1994; Castander et al.\\ 1995) suggest a significant drop in the cluster space density as compared to the value observed locally. Although the overall evidence is thus in favor of negative evolution of the cluster XLF at least for X-ray luminous clusters at redshifts well in excess of 0.3, the overall picture is anything but clear. With the completion of the BCS we are now able to provide a definitive answer to the question of whether cluster evolution is significant at low to intermediate redshifts and, in any case, provide an accurate determination of the local cluster XLF as a much-needed reference for ongoing and future evolutionary studies at higher redshifts. The implications of our findings for cosmological models of cluster evolution will be addressed in a forthcoming paper (Ebeling et al., in preparation). We assume $H_0=50$ km s$^{-1}$ Mpc$^{-1}$ and $q_0=0.5$ throughout this paper. ", + "conclusions": "Using the ROSAT Brightest Cluster Sample (BCS) as presented by Ebeling et al.\\ (1996b) we have established the local X-ray luminosity function (XLF) of clusters of galaxies within $z=0.3$ with unprecedented accuracy. We find the XLF to be well described by a Schechter function whose free parameters $A$, $L_X^{\\star}$, and $\\alpha$ we determine in a maximum-likelihood fit for all X-ray energy bands currently used within the community. Comparing our results with previous measurements of the cluster XLF we find very good agreement with the work of Piccinotti et al.\\ (1982), Henry et al.\\ (1992), and Burns et al.\\ (1996), as well as with the XLF for the B50 sample of Edge et al.\\ (1990) when the same maximum likelihood algorithm is used to determine the best Schechter function fit. We find no significant variations in the amplitude or the characteristic luminosity of the best-fitting Schechter function as a function of redshift. Also, the distribution of $V/V_{\\rm max}$ values is consistent at the 74\\% confidence level with a non-evolving space density of clusters out to $z=0.3$. Our findings do thus not confirm the claim of strong evolution at $z\\la 0.2$ made by Edge and coworkers but support the notion of Ebeling et al.\\ (1995) that the apparent signature of evolution in the B50 sample is due to a combination of its high X-ray flux limit in the $2-10$ keV band and a pronounced, if statistically insignificant, dearth of very X-ray luminous clusters around a redshift of about 0.15." + }, + "9701/astro-ph9701209_arXiv.txt": { + "abstract": "We present the highest quality \\Lya forest spectra published to date. We have complete 7.9 \\kms FWHM spectra between the \\Lya and \\Lyb emission lines of the bright, high redshift (V=15.9, $z$=3.05) QSO HS 1946+7658. The mean redshift of observed \\Lya forest clouds is $=2.7$. The spectrum has a signal to noise ratio per pixel of 2 \\kms that varies from 15 at 4190\\AA \\space to 100 at 4925\\AA. The absorption lines in the spectrum have been fit with Voigt profiles, and the distribution of Voigt parameters has been analyzed. We show that fitting Voigt profiles to high quality data does not give unique results. We have performed simulations to differentiate between true features of the line distributions and artifacts of line blending and the fitting process. We show that the distribution of H~I column densities is a power law of slope $-1.5$ from N(H~I) = $10^{14}$ \\cmm to N(H~I) $< 10^{12.1}$ \\cmm. We further show that our data is consistent with the hypothesis that this power law extends to N(H~I) = 0, because lines weaker than N(H~I) $= 10^{12.1}$ \\cmm do not have a significant H~I optical depth. At velocity dispersions between 20 and 60 \\kms the velocity dispersion ($b$) distribution is well described by a Gaussian with a mean of 23 \\kms, and a $\\sigma_b$ of 14 \\kms. Very similar N(H~I) and $b$ distributions were found at $ = 3.7$ by Lu et al. (1997), indicating no strong redshift evolution in these distributions for the \\Lya forest. However, our $b$ distribution has a lower mean and a wider dispersion than in past studies at the same redshift (eg. Hu et al. 1995) which had lower signal to noise spectra. We unambiguously see narrow \\Lya forest clouds with 14 \\kms $\\le b \\le$ 20 \\kms that cannot be accounted for by noise effects. Our data also has absorption lines with $b \\ge$ 80 \\kms that can not be explained by the blending of lower $b$ lines. We find that the lower cutoff in the $b$ distribution varies with N(H~I), from $b = 14$ \\kms at N(H~I) $= 10^{12.5}$ \\cmm to $b = 22$ \\kms at N(H~I) $= 10^{14.0}$ \\cmm. However, we see no similarly strong indication of a general correlation between $b$ and N(H~I). In contrast with previous results, we find no indication of \\Lya forest line clustering on any scale above 50 \\kms. Even among lines with $10^{13.6} < N(H~I) < 10^{14.3}$ \\cmm, which were previously thought to cluster very strongly on velocity scales between 50 and 150 \\kms, we see no clustering on any scale above 50 \\kms, although we do see a 3 $\\sigma$ clustering signal between 25 and 50 \\kms among these higher column density lines. With the distributions we have derived, we have calculated the expected He II optical depth of the \\Lya forest. If there are no lines with N(H~I) $< 10^{12.1}$ \\cmm, the \\Lya forest is unlikely to provide a significant portion of the He II optical depth at high redshift. However, if the distribution extends to N(H~I) $< 10^{9}$ \\cmm, the \\Lya forest can provide all of the observed optical depth if N(He~II)/N(H~I) $\\approx 100$. We have calculated the redshift evolution of the optical depth from the He~II \\Lya forest based upon the line distributions we have derived for the H~I \\Lya forest. If the \\Lya forest is responsible for the high redshift He~II optical depth and the spectral shape of the UV background does not change with redshift, we predict \\th $\\approx 2.4$ at $z=3.3$ to be consistent with the value of \\th previously found at $z=2.4$, provided that \\Lya forest lines with N(H~I) $<10^{13}$ \\cmm evolve like those with N(H~I) $>10^{13}$ \\cmm. ", + "introduction": " ", + "conclusions": "" + }, + "9701/astro-ph9701045_arXiv.txt": { + "abstract": "We have analyzed high dispersion and high precision spectra of 5 blue horizontal branch stars in the globular cluster M92 to establish that the projected rotational velocity for these stars ranges from 15 to 40 \\kms. This is larger than that expected based on the rotation of their main sequence progenitors, the spin down of rotation with age, and the conservation of angular momentum. Possible explanations include a rapidly rotating stellar core. An abundance analysis of these spectra of these blue HB stars in M92 yields the same results as have been obtained from the giants in this cluster. There is a hint of a trend of higher abundance as the projected surface rotational velocity increases, which could be chance and requires confirmation. \\bigskip \\bigskip \\bigskip \\bigskip \\bigskip \\bigskip \\centerline {{\\it Subject headings:} stars: rotation --- stars: chemical composition --- globular clusters: M92} \\bigskip \\bigskip \\bigskip \\bigskip \\bigskip \\bigskip \\centerline{In Press in {\\sl The Astronomical Journal.}} ", + "introduction": "There is a large suite of complicated puzzles involving the properties of the horizontal branch in globular clusters. One of these involves the distribution of stars along the horizontal branch, and how that distribution varies from cluster to cluster, i.e. the second parameter problem. (See, for example, Fusi Pecci et al 1992, 1993.) We believe that the mass loss varies from star to star along the horizontal branch in a given globular cluster, but do not understand what parameters control the mass loss near the tip of the giant branch and at the helium core flash. Another puzzle involves the possible impact of the mixing of the lighter elements, C, N, and O, and probably Na, Mg, and Al as well, produced by nuclear burning in the interior of these evolved stars near the tip of the red giant branch to their surfaces, as reviewed by Kraft (1994). We do not know how, or even if, that is related in some way to the morphology of the horizontal branch. There have also been tantalizing hints of unexpectedly large surface rotational velocities among the horizontal branch stars (Peterson 1983, 1985 a,b), but the data available have until recently been very limited. If conservation of angular momentum prevails, high surface rotation among the red giants in a globular cluster would not be expected due to their large radii, nor is it observed. There is essentially no information on the surface rotation of lower main sequence stars in globulars. High internal rotation would be impossible to detect among lower main sequence stars or among red giants, but might be revealed after mass loss at the helium flash when the stars reach the horizontal branch. It is this phenomenon which we may have detected, although other interpretations are of course possible. As discussed by Sweigart \\& Mengel (1979) and by Sweigart (1996), high rotation may influence the process and outcome of mixing, halt diffusion of elements out of the atmosphere, affect the population of the horizontal branch, etc. The advent of the Keck 10-m Telescope, and its high resolution echelle spectrograph (Vogt et al 1994), prompted us to re-examine some of these issues. ", + "conclusions": "The main sequence progenitors of the blue horizontal branch stars in M92 have a mass of about 0.85 \\Msun~ and are late F to early G dwarfs. The rotation of main sequence stars has been a subject of interest since the work of Kraft (1967). For many nearby stars, periods are now available which yield the rotation velocity itself rather than the projected velocity of rotation. These have come from extensive monitoring of the variability of the chromospheric emission in the core of the H and K lines (Baliunas \\& Vaughan 1985). Benz et al (1984) showed that in clusters as young as the Hyades, $$ is less than 12 \\kms~ for 10 stars with spectral types ranging from F8V to G5V. $v_{rot}$ as large as 50 \\kms~ is not reached in the Hyades until the spectral type is F5V, too early to be the progenitors of the blue HB stars in M92. We assume that metallicity differences do not significantly perturb this picture. In addition, it is clear from studying open clusters of varying ages younger than the Hyades that there is a spin-down of rotation with time for these stars, where the rotational decay due to magnetic braking is predicted to be of the form $t^{-0.5}$ (Skumanich 1972, Endal \\& Sofia 1981). Whether and how far this extends beyond the age of the Hyades is not known, but extrapolating blindly from 7 x 10$^8$ years (the determination of the age of the Hyades by Mazzei \\& Pigatto 1988) to ages characteristic of globular clusters, one expects a further spin-down of the rotational velocity of a factor of $\\approx$4. Based on models of horizontal branch evolution, the mean mass of the M92 blue HB stars is about 0.70 \\Msun, with a scatter of perhaps 0.03 \\Msun. (Crocker \\& Rood 1988, Lee et al 1990). We assume solid body rotation and ignore any further spin-down for the main sequence stars beyond that shown in the Hyades. We assume that the angular momentum/gm in the ${\\approx}0.15${\\Msun} lost during stellar evolution between the main sequence and the HB is the same as that of the initial star. The radii of these blue HB stars are $\\approx$3 \\Rsun, and we find that the expected rotational velocity for blue HB stars should be under 10 \\kms. Our projected rotational velocities, as well as the latest work by Peterson et al (1995), show that velocities exceeding those derived under the assumptions above are common among the blue HB stars in globular clusters. Among the many possibilities that could explain these results is that main sequence stars have a more rapidly rotating core, and as mass is lost, a higher surface rotation is seen. The evolution of differentially rotating stars is discussed by Pinsonnealt et al (1991). Peterson et al (1995) gives a detailed discussion of the myriads of other possible explanations. Several recent studies at high dispersion of field RR Lyrae stars (Clementini et al 1995, Lambert et al 1996, Peterson et al 1996, Fernley \\& Barnes 1996) demonstrate that high rotation is not seen in any of these objects, and the sample of field RR Lyrae stars observed at high dispersion is now reasonably large. The only time that line broadening can be detected is near $\\phi \\approx$ 0.85, and this presumably arises from the effect of shocks. These variables are HB stars only slightly cooler than the coolest of the blue HB stars in M92 that we examine here. We speculate this is related to the fact that in one case we are looking at field stars with a relatively young mean age, and in the other we are looking at much older objects, although it is also possible that globular cluster stars have more angular momentum ab initio than do field halo stars. This difference is quite puzzling. We have analyzed high dispersion and high precision Keck/HIRES spectra of 5 blue horizontal branch stars in the globular cluster M92 to establish that the projected rotational velocity for these stars ranges from 15 to 40 \\kms. This is larger than that expected based on the assumed rotation of their main sequence progenitors, the spin down of rotation with age, and the conservation of angular momentum. Possible explanations include a rapidly rotating stellar core. An abundance analysis of these spectra of blue HB stars in M92 yields the same results obtained from the giants in this cluster. There is evidence of a small increase in abundance of the heavier elements as $v_{rot}(sin(i))$ increases, but the sample urgently needs to be enlarged. We now feel confident that our methods are valid and that we are ready to proceed to analyze the spectra of HB stars in much more metal-rich globular clusters." + }, + "9701/astro-ph9701103_arXiv.txt": { + "abstract": "We investigate the importance of projection effects in the identification of galaxy clusters in 2D galaxy maps and their effect on the estimation of cluster velocity dispersions. From large N-body simulations of a standard cold dark matter universe, we construct volume-limited galaxy catalogues that have similar low-order clustering properties to those of the observed galaxy distribution. We then select clusters using criteria tailored to match those employed in the construction of real cluster catalogues such as Abell's. We find that our mock Abell cluster catalogues are heavily contaminated and incomplete. Over one third (34$\\pm$6 per cent) of clusters of richness class R$\\geq$1 are miclassifications arising from the projection of one or more sub-clumps onto an intrinsically poor cluster. Conversely, 32$\\pm$5 per cent of intrinsically rich clusters are missed altogether from the R$\\geq$1 catalogues, mostly because of statistical fluctuations in the background count. Selection by X-ray luminosity rather than optical richness reduces, but does not completely eliminate, these problems. Contamination by unvirialised sub-clumps near a cluster leads to an overestimation of the cluster velocity dispersion which can be very substantial even if the analysis is restricted only to galaxies close to the cluster centre. Thus, the distribution of cluster masses -- often used to test cosmological models -- is a highly unreliable statistic. The median value of the distribution, however, is considerably more robust because the main effect of contamination is to create an artificial tail of high velocity dispersion clusters. Improved estimates of the cluster velocity dispersion distribution require constructing new cluster catalogues in which clusters are defined according to the number of galaxies within a radius about three times smaller than the Abell radius. ", + "introduction": "Clusters of galaxies are a major source of cosmological information. Because of their large luminosity they can be detected, and their properties can be measured with relative ease, out to large distances. This makes it possible to exploit their special characteristics as the most massive nonlinear objects in the Universe. In hierarchical clustering theories for the formation of structure, clusters are associated with the rare high peaks in the primordial density field on scales of a few megaparsecs. As a result, their mass and abundance are very sensitive to the amplitude of mass fluctuations on these scales (Frenk \\etal 1990; White, Efstathiou \\& Frenk 1993; Viana \\& Liddle 1996; Eke, Cole \\& Frenk 1996b). The epoch of cluster formation and the rate at which the cluster population builds up is, similarly, a strong function of the mean density parameter, $\\Omega$ \\cite{laco93,eke96b}, as is their degree of internal substructure (Richstone, Loeb \\& Turner 1992; Mohr \\etal 1995; Wilson, Cole \\& Frenk 1996). The clustering properties of clusters depend primarily on the shape of the power spectrum of mass fluctuations and have been a subject of much debate for over 20 years \\cite{bs83,dalton94,eke96a}. Finally, rich clusters have recently been used to map the local density field (Plionis \\etal 1996, in preparation). The use of galaxy clusters as cosmological diagnostics relies on the availability of statistical samples, selected according to a well-defined property, such as richness, mass, or X-ray temperature. Traditionally, the source of such samples has been Abell's (1958) cluster catalogue. The integrity of the Abell catalogue, however, has been questioned over the years (e.g. Fesenko 1979a,b; Lucey 1983; Frenk \\etal 1990). Even Abell himself made it quite clear that the completeness and homogeneity of his catalogue were suspect. Partly to overcome these shortcomings new cluster catalogues were constructed in the early 1990s based, as Abell's, on photographic material, but replacing eye-ball identifications by automated scans of digitized plates. These procedures have produced the APM \\cite{dalton92} and Edinburgh-Durham cluster catalogues (EDCC; Lumsden \\etal 1992). Computer manipulation of the galaxy images has allowed a degree of uniformity and repeatability to be reached that was impossible in Abell's days. Nevertheless, Abell clusters remain the best studied, and Abell's catalogue the main source from which samples are drawn for statistical studies. Whether by eye or by computer, catalogued clusters are identified as two-dimensional objects, seen against a strongly clustered background. The enormous column depth to a cluster makes projection effects inevitable. Indeed, spectroscopic follow-up of Abell and ACO clusters \\cite{aco} often reveals several clumps of galaxies lined up in the direction to a rich cluster (see Katgert \\etal 1996 for a recent study of a large sample). Although these clumps enhance the apparent richness of the cluster, the dominant concentration along the line-of-sight is often rich enough to emit X-rays (e.g. Briel \\& Henry 1993). This fact alone, however, says little about the importance of projection effects or the completeness of optically selected cluster samples. Since the abundance of clusters declines very rapidly with richness or mass, even a small amount of contamination can compromise statistical studies in which completeness and/or homogeneity are required. The cluster two-point correlation function is a good example. Values for the correlation length, $r_0$, differing by almost a factor of two have been strongly advocated by various workers \\cite{bs83,postm92,efst92,nich92,dalton94}. According to Bahcall \\& West (1992) the differences are due to the different limiting richnesses of the various samples but others have claimed that they are due (at least in part) to a misclassification of poor clusters which are placed into a higher Abell richness class as a result of contamination by the halos of rich clusters \\cite{suth88,dekel89,efst92}. Eke \\etal \\shortcite{eke96a} have argued that a combination of different selection procedures and a richness dependence of the clustering strength also contributes to these differences. While even distant sub-clumps artificially enhance the apparent richness of a cluster, subclustering in its immediate vicinity causes its velocity dispersion to be overestimated (e.g. Frenk \\etal 1990). This effect is the likely cause of the poor correlation between the X-ray temperature and velocity dispersion of the most massive clusters \\cite{dajofo95} and it vitiates comparisons between cluster masses determined from X-ray, optical and gravitational lensing data \\cite{fahl94}. The distribution of cluster masses or velocity dispersions has been used as a discriminant of different cosmological models (e.g. Weinberg \\& Cole 1992; Bahcall \\& Cen 1993; Lubin \\etal 1996). These comparisons tend to rely heavily on the behaviour of the high mass end of the distribution which, unfortunately, is particularly sensitive to contamination due to substructure. Masses derived from X-ray data are more reliable (Evrard, Metzler \\& Navarro 1996; but see Balland \\& Blanchard 1995), but since cluster samples are almost invariably selected by their optical properties, the inferred distributions of X-ray properties are also subject to the kind of uncertainties discussed above. Furthermore, by virtue of the fact that optically selected cluster catalogues all have a lower cut-off in richness, there is an in-built bias in these samples against low mass clusters. It seems clear that the use of cluster properties as cosmological diagnostics requires a detailed understanding of the biases introduced by projection effects and contamination in cluster catalogues. The aim of this paper is to set up a methodology for quantifying such biases using mock galaxy catalogues constructed from N-body simulations. In this paper we analyze mock catalogues constructed from standard $\\Omega=1$ cold dark matter (CDM) simulations and we concentrate on Abell clusters. Our procedures, however, can readily be extended to other cosmologies and other cluster catalogues. We also use our mock catalogues for testing alternative procedures for defining clusters and estimating their properties which avoid some of the biases present in Abell's catalogue. Our work extends previous analyses by Frenk \\etal\\shortcite{fwed} and White (1991,1992) who used similar techniques. An earlier assessment of the completeness and contamination of Abell's catalogue, using Monte-Carlo simulations, was carried out by Lucey \\shortcite{lucey83}. He found that between 15 and 25 per cent of rich Abell clusters have a true membership that is less than half the number observed. However, his models contained no dynamical information, and they did not take into account the clustering properties of galaxies and clusters. In Section~2 we give technical details of our simulations and our method for constructing galaxy catalogues from which mock Abell catalogues are derived. The galaxy catalogues encapsulate the essential characteristics of the real situation, although we have not tried to reproduce the observed properties in every detail. Section~3 presents an analysis of the completeness of the cluster catalogues. By identifying groups along the line-of-sight to each cluster, we determine the properties of the main concentration of galaxies, and those of clumps projected onto the cluster. The various categories of contamination we identify are described in Section~4. In Section~5 we discuss how projection effects influence estimates of cluster velocity dispersions derived from radial velocity measurements and illustrate how careful cluster selection and interloper removal can improve upon the accuracy of these estimates. In Sectio~6 we apply a popular statistical test for substructure to our data and demonstrate its potential for flagging clusters whose velocity dispersion estimates are strongly affected by substructure. In Section~7 we present the cluster-cluster correlation function, and illustrate how overlapping clusters can affect the amplitude of this function. Finally, we present a summary of our main results in Section~8. \\begin{figure} \\centering \\centerline{\\epsfysize=9.5truecm \\figinsert{fig01.eps}{Figure 1}} \\caption{The real space two-point correlation function of galaxies in our catalogues compared with observations. The solid circles show $\\xi(r)$ obtained by averaging over 8 independent simulated galaxy catalogues. The open circles show Baugh's (1996) deprojection of the angular correlation function of the APM galaxy catalogue. The triangles show the real-space correlation function derived by Loveday \\etal (1995) from the Stromlo-APM galaxy redshift survey.} \\end{figure} ", + "conclusions": "We have investigated two sources of uncertainty in the construction and analysis of cluster catalogues selected from 2D galaxy maps. The first is the influence of projection effects on the identification and richness classification of clusters. The second is the influence of local subclustering on estimates of cluster velocity dispersions. We have found that both these effects are large and compromise the use of cluster properties as cosmological diagnostics. Our conclusions are based on the analysis of mock cluster catalogues. Although the clusters themselves are identified in 2D galaxy maps using criteria patterned on those employed in real catalogues, in the simulations we have access to full 3D spatial and velocity information for the model galaxies. These were identified using the high peak model of biased galaxy formation (the peak-background split technique) in an $\\Omega=1$ standard CDM universe with primordial power spectrum normalized so as to obtain the correct abundance of rich Abell clusters. The resulting galaxy autocorrelation function is similar to the observed one in the relevant range of separations $21$ spectrum of metric perturbations. The model contains a symmetry breaking field that tunnels to its true vacuum, producing a single bubble inside which hybrid inflation drives the universe to almost flatness. In order to obtain density perturbations with $n > 1$ we analyse a recently proposed new version of hybrid inflation scenario called tilted hybrid inflation. In this scenario, unlike in the previously known versions of hybrid inflation, a considerable tilt of the spectrum can be obtained without fine-tuning. The stage of inflation in this model is rather short, which allows us to obtain an inflationary universe with $\\Omega < 1$ in a more natural way. We study the separate contribution of scalar perturbations coming from the continuum subcurvature modes, the discrete supercurvature mode and the bubble wall mode to the angular power spectrum of temperature fluctuations in open inflation. We derive bounds on the parameters of the model so that the predicted spectrum is compatible with the observed anisotropy of the microwave background. ", + "introduction": "Until recently, inflation was associated unequivocally with a flat universe, due to its tendency to drive the spatial curvature so effectively to zero. However, it is now understood that inflation comprises a wider class of models, some of which may give rise to an open universe at present~\\cite{open,BGT,LM}. Such models generically contain a field trapped in a false vacuum, which tunnels to its true vacuum via the nucleation of a single bubble inside which a second period of inflation drives the universe to almost flatness. These models recently became rather popular because in an open universe it is possible to reconcile a large value of the Hubble constant~\\cite{HST} with the large estimated age of globular clusters~\\cite{GC}. But even independently of this issue, open inflation models have several interesting features that single them out from other cosmological models. For a long time it seemed impossible to make any physical sense of the (mathematically consistent) open universe model because it presumed that an infinite universe appeared from nowhere at a single moment of time, being perfectly synchronized over an infinitely large length scale. This leads to the horizon problem in its strongest form. Also, whereas it seemed possible that a small closed universe could be created by `tunneling from nothing' and then experience a period of inflation, a similar event for an infinite open universe seemed impossible. Indeed, one may argue that such processes should be suppressed by $e^{-|S|}$, where $S$ is the action of the instanton describing the universe creation \\cite{LVIL}. Therefore one could expect that quantum creation of an infinitely large open universe should be forbidden because this would involve tunneling with infinite action. The models of Refs.~\\cite{open,BGT,LM} show for the first time that a consistent physical model of an infinite homogeneous universe is possible, and that such a universe can appear as a result of quantum tunneling. The bubble containing an open universe looks finite from the point of view of an outside observer, and the probability of its creation is finite. Meanwhile an observer inside the bubble would see an infinite open universe. Another distinct feature of such models is that the homogeneity problem is solved not through the exponential expansion, as in the usual inflationary models, but through the exponential suppression of the probability of production of spherically asymmetric bubbles~\\cite{LM}. This way one solves the homogeneity problem independently from the flatness problem. The origin of structure is still related to amplified quantum fluctuations of the field that drives inflation inside the bubble~\\cite{sasaki,BGT}. However, in the spectrum of temperature fluctuations there appears a new set of discrete supercurvature modes~\\cite{LW,GZ}, associated with the open de Sitter vacuum~\\cite{sasaki,YST} and the bubble wall fluctuations at tunneling~\\cite{LM,Hamazaki,bubble,Garriga}, which could be made small in some of the models~\\cite{super}. There is some evidence that the observations made in a wide range of scales, from horizon size to large clusters of galaxies, constrain open inflation models (with small $\\Omega_0\\sim0.3$--$0.4$) to have a `tilted' spectrum of density perturbations with the spectral index $n>1$ ~\\cite{WS}, and essentially no other contribution, either from gravitational waves or supercurvature modes. The open inflation models considered so far predict a tilted $n<1$ spectrum, see e.g. Ref.~\\cite{induced}, which may contradict observational data for models with very small $\\Omega_0$~\\cite{WS}. In this paper, we consider a model of open inflation in which the second stage inside the bubble is driven by hybrid inflation~\\cite{hybrid}. Models of hybrid inflation are known to provide a blue tilted spectrum of metric perturbations, with negligible contribution of gravitational waves~\\cite{LL93,CLLSW,GBW,Lyth}. However, in most of the hybrid inflation models developed so far the tilt is extremely small; in order to achieve a considerable tilt of the spectrum in hybrid inflation one needs to fine tune the parameters of the model~\\cite{GBW}. In what follows we will consider a new class of hybrid inflation models where a significant tilt can be achieved in a natural way, which we call {\\it tilted hybrid inflation}, see Ref.~\\cite{GBL}. We will make tilted hybrid inflation a part of an open inflation scenario. To make sure that supercurvature and bubble wall modes do not distort the CMB too much, we compute the angular power spectrum of temperature anisotropies for both the continuum modes and the supercurvature modes and find that there are models in which all constraints are satisfied. Apart from scalar metric perturbations, open inflation also produces a primordial spectrum of gravitational waves, whose amplitude and scale dependence in single-bubble open inflation models has only recently been known~\\cite{TS}. In order to compare with observations there still remains to be computed the corresponding angular power spectrum of CMB temperature fluctuations. We will consider the constraints on the parameters of tilted hybrid models from such a tensor component of the CMB anisotropies in a future publication. ", + "conclusions": "In this paper we have studied an open inflation model in which the second period of expansion inside the bubble is driven by tilted hybrid inflation~\\cite{GBL}. This model has a characteristic signature of a tilted $n>1$ spectrum. Such a spectrum is nowadays the best candidate for explaining the CMB anisotropies in case the universe is wide open, with $\\Omega_0 < 0.4$~\\cite{WS}. We have computed the angular power spectrum of temperature fluctuations in the CMB induced by the continuum of scalar subcurvature modes, as well as the de Sitter supercurvature mode and the bubble wall mode. We have found a set of constraints that the parameters of such a model should satisfy in order to agree with observations. The parameters of our model satisfy all the constraints in a natural way. The contribution of a primordial spectrum of gravitational waves to the complete angular power spectrum of temperature anisotropies and the associated constraints on the parameters of the open hybrid model will be discussed in a separate publication. Note that our model requires the existence of three scalar fields, which is a long shot from the simplest inflationary models describing only one scalar field. On the other hand, our model represents a simple combination of the hybrid inflation scenario and the theory of a scalar field $\\sigma$ with a metastable vacuum state $\\sigma=0$. The hybrid inflation scenario recently became very popular. It allows inflation to occur at sub-Planckian values of scalar fields, which makes it easier to realize it in the context of supersymmetric models. The scalar field $\\sigma$ with a metastable vacuum state is necessary for open inflation anyway. Thus our model consists of only two building blocks, each of which is rather natural in the context of inflationary cosmology. There are many ways to combine these two models together, so the model described in this paper is not unique, and it is quite possible that a simpler and more elegant model can be proposed. Our main purpose here was simply to show that open inflation models with blue spectra do exist. On the other hand, the tilted hybrid inflation which we used in our scenario has a naturally short duration of inflation after the tunneling with the open universe formation. This is exactly what we need in order to have an inflationary universe with $\\Omega<1$. In principle, one can obtain a tilted spectrum of perturbations in an entirely different way, by combining the open inflation model with the recently proposed model of Ref.~\\cite{LinMukh}, in which one may have strongly tilted isocurvature or adiabatic non-Gaussian perturbations. It might be possible to obtain a model with all desirable properties that will require only two scalar fields. To avoid misunderstandings, we should emphasize that we sincerely hope that at the end of the day observational data will show that the universe is flat, and that perturbations of the metric are scale-invariant, as suggested by the simplest inflationary models. However, when the universe was built we have not been consulted. Our final goal is to propose an internally consistent theory that will correctly describe observational data. We are very encouraged that a consistent theory of an open universe does exist. So far this theory has been formulated only in the context of inflationary cosmology. This means that the consistent cosmological models describing large nearly homogeneous universe with $\\Omega < 1$ do not contradict inflation, and, moreover, they have so far been constructed only within the context of inflation. In this paper we have shown that it is not difficult to modify the spectrum of perturbations of the metric in such models, which may be necessary to make them consistent with observations." + }, + "9701/astro-ph9701167_arXiv.txt": { + "abstract": "The data of {\\it BVR} observations of the middle-age radio pulsar PSR 0656+14 on January, 20/21 at the BTA (6-m) are presented. The brightness is determined in Cousins {\\it B} filter $B\\approx 25.1$ with $\\lambda_{eff}=4448 \\rm \\AA$ in adjacent for HST F130LP long-pass filter of a star-like object, coinciding with the position of VLA radio source. Relatively large observed {\\it V} and {\\it R} fluxes ($\\lea 3\\sigma\\ or\\ > 10^{-30}\\ ergs\\ cm^{-2}\\ s^{-2}\\ Hz^{-1}$) can witness a non-thermal nature of optical radiation of this pulsar up to $\\lambda \\approx 6600\\ \\rm \\AA$. Most probably, in the UV-optical ({\\it BVR}) spectral range a power-law spectrum is superimposed on the thermal-like radiation of the entire neutron star surface what can be related to a mechanism itself of the pulsar activity. \\keywords {\\mbox{ pulsars, ground-based observations, CCD photometry}} ", + "introduction": "At present above 700 radio pulsars are discovered, but the optical radiation can be assumed reliably detected only for some of them: the famous Crab pulsar, PSR 0540-69, PSR 1509-58, the Vela pulsar, the gamma-ray + X-ray pulsar Geminga, and for PSR 0656+14 (\\cite{12}; \\cite{4}; \\cite{13}; \\cite{9}; \\cite{3}; \\cite{5}; \\cite{11}). I.e. the probable optical companions are already detected the which behavior can be investigated by more careful study of their spectra and temporal variability both in X-ray, gamma, radio, and optical. Recent observations at the Hubble Space Telescope (HST) in UV-optical range led to the detection of corresponding probable UV-optical counterparts of another two isolated pulsars PSR 0950+08 and PSR 1929+10 (\\cite{11}). PSR 0656+14 was also identified recently in optical (\\cite{4}) and immediately after that the study of its spectrum was started at the HST (\\cite{11}). At the 6-meter telescope (BTA) this pulsar and others is studied within the framework of a program of wide-band ground-based photometry of nearest pulsars of the Northern sky. For a middle-aged pulsar PSR 0656+14 ($\\tau =\\frac{P}{2\\dot{P}} = 110000$ yr) we suppose to fulfill a multi-color photometry for the purpose of refinement of the nature of optical radiation of the isolated neutron star (INS), by supplementing HST UV-optical observations with ground-based {\\it BVRI}-observations. Certainly, the basic (global) goal pursued by many groups at the investigation of INSs-pulsars is to select a thermal component of radiation (including the optical one) for the pulsars of the age of $>~10^5$ years arising from the entire neutron star surface. This problem is still actual since as is noted in all recent papers on INSs it would allow us (together with the study in {\\it EUV} and X-ray ranges) to refine the thermal evolution of these compact objects and to approach in the end the correct equation of state of matter for their interior regions with supernuclear densities. A possibility is now actively discussed of a presence in the deep interior of \"neutron\" stars of a pion or quark condensate , a superfluidity, and others. See, for example, Umeda, Tsuruta, and Nomoto (1994); Meyer, Pavlov, and Meszaros (1994). Though, it should be remarked that UV-optical thermal radiation was apparently observed only for two active and more old pulsars: PSR 1929+10 and (probably) PSR 0950+08 (\\cite{11}). In other cases the thermal radiation from INSs is basically observed in {\\it EUV} and soft X-ray bands. The study of spectra of the nearest isolated pulsars, including the X-ray brightest PSR 0656+14, was begun with the study in X-ray (\\cite{14}). In particular, now for PSR 0656+14 the observation in soft X-ray range of the ROSAT observatory (\\cite{7}) are also used. The high quality of X-ray spectra allows to determine rather precisely an effective temperature of the radiating surface of a neutron star by means of \"fitting\" the observed spectra to fit the black-body like radiation in Wien region. However, the most recent observations of isolated pulsars, such as Geminga and PSR 0656+14, including optical investigation (\\cite{3}; \\cite{11}) showed that in optics the effects can become essential which are related either to the presence of geometrically thin ($\\sim 1.5$ cm), but optically thick atmosphere, or with the influence of magnetic field , or with the non-thermal contribution of radiation from polar caps, or with some other non-thermal effects. One way or another, it turns out that a simple black-body fitting an observed spectrum can be quite non-adequate to not only X-ray range (\\cite{10}), but especially to optical range (\\cite{3}; \\cite{11}). I.e. in the UV-optical range a non-thermal radiation could dominate the optical spectrum at least in the middle-aged pulsars. Though the study of non-thermal radiation is interesting by itself from the point of view of elucidation of physical conditions in pulsar magnetospheres and refinement of a theory of pulsar emission, but it is nevertheless a \"barrier\" in the movement to the basic goal - the elucidation of the main question: \"What the interior of neutron stars consist of ?\" On the other side, hopefully, the presence of a just non-thermal component of radiation can increase considerably the luminosity of these objects in optical. The last was also directly confirmed, in particular, by our BTA observations of PSR 0656+14 in {\\it B, V, R} filters, about what the 2-d section of this paper says. We obtained for the first time the estimations of the brightness in this bands from the Earth. Though in {\\it V} filter the observational material of approximately the same quality was already obtained at two ESO telescopes by \\cite{4}. The first attempt to observe PSR 0656+14 in optical was undertaken by Cordova {\\it et al.} (1989) after the identification of this pulsar in X-rays. Though an optical counterpart was not then detected, this observation showed that the corresponding region around the VLA position of the pulsar is not espacially crowded and does not contain too bright objects nearer $\\sim 5$ arcsec from the pulsar. Thereupon a successful ground-based optical observation of PSR 0656+14 which is the X-ray brightest from all \"normal\" radio pulsars was fulfilled in 1989 at the 3.6-m ESO telescope for a total exposure time of 60 minutes in {\\it V} filter and in 1991 with the NTT for the 70 minutes total exposure in {\\it V} bandpass (\\cite{4}). In both cases the authors detected an object which coincides well with VLA position of PSR 0656+14 as measured by Thompson and Cordova (1994). The corresponding object has $V \\sim 25$ with the error of 0.5 mag, what corresponds to the 3$\\sigma$ level of detection. However, a large stellar magnitude of optical counterpart together with uncertainties in the estimate of distance to the pulsar (100-700 pc) makes difficult the classification of optical both thermal and non-thermal radiation, especially in terms of uncertain interpretation of X-ray data (\\cite{7}). New observations of this pulsar were carried out at the HST with the UV-sensitive Faint Object Camera by Pavlov {\\it et al.,} (1996). The observations of the pulsar candidate for PSR 0656+14 were fulfilled in F130LP filter with the band width $\\lambda\\lambda = 2310 -4530 \\rm \\AA$, with center at the $3365 \\rm \\AA$, including the radiation in standard {\\it U} and {\\it B} filters. The exposure was 4755 sec. Near the VLA position of PSR 0656+14 in the deep $7\\farcs4\\times7\\farcs4$ images there is only one point-like object with $m_{130LP} = 25.19\\pm 0.04$ ($S/N=52$). Results of first observations at the 6-m telescope in more narrow spectral bands for the candidate identification for PSR 0656+14, supplementing the observations of the group of Pavlov {\\it et al.} (1996) in space are presented in Section 2. Thus, the goals of this paper are: 1). to confirm by our observations a non-thermal nature of optical radiation of the candidate identification for PSR 0656+14; 2). to understand also how the optical spectrum of this middle-age pulsar differs from an analogous spectrum of Geminga (\\cite{3}) which is much alike to it - another middle-age INS. To draw a power low spectrum Pavlov {\\it et al.} (1996) used the point of $V\\approx 25$ obtained in observations at NTT by Caraveo {\\it et.al.} 1994) and their own point obtained at HST in F130LP filter. On the one hand, the measurements of brightness of the optical candidate identification for PSR 0656+14, which has been gained in \"grond-based\" {\\it B} band, confirmed indeed a non-thermal nature of optical radiation of this pulsar in optical. On the other hand, {\\it BVR} observations at the 6-m telescope give an opportunity to say more reliably about a power-law spectrum in optical, but not about a cyclotron feature on the thermal continuum, which is the Geminga case. Section 3 describes that in more details. The conclusion notes that in all cases of observations at the 4 telescopes and in different optical spectral bands the PSR 0656+14 turns out to be 1.5-2 stellar magnitude brighter than was expected before that on the basic of simple black-body fitting of X-ray and optical data. ", + "conclusions": "So, from the result of observation at 4-th (3.6-m ESO, NTT, HST, BTA) telescopes in different optical spectral bands the PSR 0656+14 candidate turns out to be 1.5-2 stellar magnitude brighter than was expected earlier, proceeding from a simple black-body fit of X-ray + optical data. The radiation of PSR 0656+14 is basically non-thermal indeed in optical, though the entire neutron star surface radiation can become dominating in the far UV range. As to our {\\it B} estimate of brightness, together with HST F130LP and ground-based observations by Caraveo {\\it et al.,}(1994) in {\\it V} filter, the non-thermal spectrum can be approximated by one of power lows suggested in the paper by Pavlov {\\it et al.,} (1996). Most probably this object has no sharp decrease of flux in the more narrow (in comparison to F130LP) \"ground based\" {\\it B} band as compared to the flux in {\\it V} and {\\it R} filters, which (a dip) is apparently observed from Geminga (\\cite{3}). To confirm or to rule out the tendency of continuation of the PSR 0656+14 optical candidate spectrum into the red range with one or another power-law and also to have reliable data for development of quantitative models of non-thermal radiation from INSs it is necessary to carry out additional {\\it BVRI} observations. Further multicolor photometry would be needed. In particular, the Crab- pulsar spectrum is studied rather in details in this range and it is flat indeed with $\\alpha = -0.11\\pm0.13$ (\\cite{12}), though all other features of these two pulsars are much different. The authors sincerely thank George Pavlov for the most fresh information on the HST observations of INSs and for sending us the paper of Pavlov {\\it et al.,} (1996) prior to publication. We thank also Alexander Kopylov for the help in observations and active discussions of the obtained results and Tatyana Sokolova for preparing this text to publication." + }, + "9701/astro-ph9701217_arXiv.txt": { + "abstract": "\\noindent {\\normalsize Hot (explosive) hydrogen burning or the Rapid Proton Capture Process (rp-process) occurs in a number of astrophysical environments. Novae and X-ray bursts are the most prominent ones, but accretion disks around black holes and other sites are candidates as well. The expensive and often multidimensional hydro calculations for such events require an accurate prediction of the thermonuclear energy generation, while avoiding full nucleosynthesis network calculations. In the present investigation we present an approximation scheme which leads to an accuracy of more than 15 per cent for the energy generation in hot hydrogen burning from $10^8$--$1.5\\times 10^9$ K, which covers the whole range of all presently known astrophysical sites. It is based on the concept of slowly varying hydrogen and helium abundances and assumes a kind of local steady flow by requiring that all reactions entering and leaving a nucleus add up to a zero flux. This scheme can adapt itself automatically and covers situations at low temperatures, characterized by a steady flow of reactions, as well as high temperature regimes where a $(p,\\gamma)$--$(\\gamma,p)$-equilibrium is established, while $\\beta^{+}$-decays or $(\\alpha,p)$-reactions feed the population of the next isotonic line of nuclei. In addition to a gain of a factor of 15 in computational speed over a full network calculation, and an energy generation accurate to more than 15 per cent, this scheme also allows to predict correctly individual isotopic abundances. Thus, it delivers all features of a full network at a highly reduced cost and can easily be implemented in hydro calculations. \\vspace{0.5 cm} \\noindent {\\em Subject headings\\/}: nuclear reactions, nucleosynthesis, abundances --- stars: novae --- X-rays: bursts \\/} ", + "introduction": "Close binary stellar systems can exchange mass, when at least one of the stars fills its Roche Lobe. After a critical mass $\\Delta M$ of unburned transferred matter is accumulated on the surface of the accreting star, ignition sets in, typically under degenerate conditions when the accreting object is a white dwarf or neutron star. Degenerate conditions for which the pressure is not a function of temperature, prevent temperature adjustment via pressure increase and expansion and cause a thermonuclear runaway and explosive burning. For white dwarfs, a layer of $10^{-5}$--$10^{-4}$ M$_\\odot$ forms, before pycnonuclear ignition of hydrogen burning sets in (e.g. \\shortciteNP{sugimoto80}; \\shortciteNP{starrfield93}; \\shortciteNP{coc95} ) and causes a nova event (explosive H-burning on white dwarfs) with maximum temperatures of $\\approx$(2--3)$\\times 10^8$ K and a total energy release of $10^{46}$--$10^{47}$ ergs. The burning takes 100--1000 s before the partially burned hydrogen envelope is ejected. X-ray bursts (for an observational overview see \\citeANP{lewin93} 1993) involve accreting neutron stars with an unburned, hydrogen-rich surface layer. The critical size of the hydrogen layer before ignition is as small as $10^{-12}$ M$_\\odot$. Temperatures of the order (1-2)$\\times 10^9$ K and densities $\\rho$$\\approx$$10^{6}$--$10^{7}$ g cm$^{-3}$ are attained (see e.g. \\shortciteNP{wallace81}; \\shortciteANP{ayasli82} 1982; \\citeANP{hanawa83} 1983; \\citeANP{woosley86} 1986; \\shortciteANP{taam93} 1993; \\citeANP{taam96} 1996). This explosive burning with rise times of about 1--10 s leads to the release of $10^{39}$--$10^{40}$ ergs. Many of the observed features and general characteristics are understood, however, there is still a lack of a quantitative understanding of the detailed observational data. Another aspect is that the explosion energies are smaller than the gravitational binding energy of the accreted hydrogen envelope. An evenly distributed explosion energy would not unbind and eject matter. It remains to be seen whether an interesting amount of matter can escape the neutron star. Super-Eddington X-ray bursts (\\shortciteNP{taam96}) are the best candidates for this behavior. A description of hot (explosive) hydrogen burning has been given by \\shortciteANP{arnould80} (1980), \\shortciteANP{wallace81} (1981), \\shortciteANP{ayasli82} (1982), \\shortciteANP{hanawa83} (1983), \\shortciteANP{wiescher86} (1986), \\shortciteANP{wormer94} (1994), \\shortciteANP{nuc_phys_thie} (1994) [see also \\shortciteANP{biehle91} (1991), \\shortciteANP{biehle94} (1994), \\citeANP{cannon92} (1992), and \\shortciteANP{cannon93} (1993) for Thorne-Zytkow objects, but see the recent results by \\citeANP{fryer96} (1996) which puts doubts on their existence]. The burning is described by proton captures, $\\beta^+$-decays, and possibly alpha-induced reactions on unstable proton-rich nuclei, usually referred to as the rp-process (Rapid Proton Capture Process). Cross sections can either be obtained from the best available application of present experimental knowledge, e.g. the determination of resonance properties from mirror nuclei and transfer and/or charge exchange reaction studies, or actual cross section measurements, like for $^{13}$N($p,\\gamma)^{14}$O, the first reaction cross section analysed with a radioactive ion beam facility (see \\shortciteANP{champagne92} 1992 and references therein). There exist two major motivations in nuclear astrophysics (a), to understand the reaction flow to a necessary degree, in order to predict the correct energy generation required for hydrodynamic, astrophysical studies, and (b), to predict a detailed isotopic composition which helps to understand the contribution of the process in question to nucleosynthesis in general. Our main motivation in this paper is (a), to provide a fast energy generation network as a tool for nova and X-ray burst studies. This might be underlined by the fact that novae and X-ray bursts seem not to be major contributors to nucleosynthesis due to the small ejected masses involved or the question whether gravitational binding can be overcome at all. [However, novae can be important for nuclei like $^{7}$Li, $^{15}$N, $^{22}$Na, $^{26}$Al, and even Si and S; and super-Eddington X-ray bursts (\\shortciteANP{taam96} 1996) will be able to eject some matter, probably containing some light p-process elements.] It will turn out at the end that our efficient approximation and energy generation scheme can also be used to predict abundances accurately. The nucleosynthesis for nuclei above Kr will, however, be discussed in a second paper (\\shortciteANP{schatz96} et al. 1996). In order to understand the energy generation correctly, we have to be able to understand the main reaction fluxes. In the past we performed a series of rp-process studies (\\shortciteANP{wiescher86} 1986 to Ar, and \\shortciteANP{wormer94} 1994, extending up to Kr). The reaction rate predictions were based on resonance and direct capture contributions for proton-rich nuclei, making use of the most recent experimental data. Several ($p,\\gamma$)-reaction rates below mass A=44 have been recalculated by \\shortciteANP{herndl95} (1995) in the framework of a shell model description for the level structure of the compound nucleus. A preliminary analysis of the major aspects was given by \\shortciteANP{nuc_phys_thie} (1994). In the following section \\ref{reacflow} we will present this in more detail and discuss the constraints which an approximation scheme has to fulfill in the whole range of temperatures occurring in explosive hydrogen burning environments, i.e. $10^8$--$1.5\\times 10^9$ K. The approximation scheme will be presented in section \\ref{approxscheme}, its application and comparison with full network calculations in section \\ref{results}, followed by concluding remarks in section \\ref{conc}. ", + "conclusions": "\\label{conc} The main motivation for the present investigation was to find a fast and efficient approximation scheme which permits to predict the energy generation in explosive hydrogen burning for a large range of conditions, i.e. temperatures $(0.2$$\\le$$T_{9}$$\\le$$ 1.5)$ and densities $(10^{4}$~g~cm$^{-3}$$\\le$$\\rho$$\\le 10^{6}$~g~cm$^{-3}$) respectively, for applications in hydro calculations which cannot afford full nuclear network. This goal has been achieved. The energy generation of full network calculations (150 nuclei) could be reproduced with a high accuracy, resulting in deviations of 5 to a maximum of 15 per cent while gaining a factor of about 15 in computational speed. The approximation scheme discussed in this paper is therefore well suited for realistic hydro calculations of novae or X-ray bursts and other possible sites of explosive hydrogen burning. The additional advantage of this method is that it also permits to predict isotopic abundances with a similar accuracy and can thus even replace full network calculations for this purpose (see Figure \\ref{fig ar33615}). \\resetfg \\begin{figure}[tbn] \\epsscale{0.75} \\figurenum{11} \\plotfiddle{ar33615p1.ps}{6 cm}{0}{30}{30}{-100}{-20} \\figcaption[ar33615p1.ps]{Comparison of the $^{33}$Ar-abundances calculated by a full network calculation and the approximation scheme. \\label{fig ar33615}} \\end{figure} We were also able to determine key nuclear properties, which directly enter the precision of calculations for explosive hydrogen burning, its energy generation and abundance determination. We have identified the even-$Z$ T$_z$=--1/2 nuclei like $^{23}$Mg, $^{27}$Si, $^{31}$S, $^{35}$Ar and $^{39}$Ca, as essential targets for proton captures which are in competition with $\\beta$-decays. As their small reaction Q-values (less than 2~MeV) do not permit the application of statistical model cross sections, they have to be determined experimentally. Thus, these results are an ideal guidance for future experimental investigations with radioactive ion beams. Nuclei closer to stability permit the application of statistical model cross sections. Nuclei more proton-rich than $T_z$=-1 are only populated for conditions when a $(p,\\gamma)$--$(\\gamma,p)$-equilibrium is established. Thus, only their masses, spins and half-lives enter the calculation. Connecting $(\\alpha,p)$-reactions have high enough Q-values, even close to the drip-line, that the level densities are sufficient for statistical model cross sections (this is true at least up to Ca and Ti, where alpha-induced reactions can play an important role). \\nopagebreak Therefore, the remaining experimental properties which should be determined and enter the calculations in a crucial way are $\\beta^+$-decay half-lives and nuclear masses. The latter control the abundances in a $(p,\\gamma)$--$(\\gamma,p)$-equilibrium. Beyond Se, the details of the proton drip-line are actually not that well known yet, but can influence the endpoint of the rp-process. First calculations which include 2p-capture reactions (Schatz et al. 1996) indicate that it seems possible to produce nuclei with A=90-100 in the rp-process under X-ray burst conditions, depending on the choice of mass formulae and half-life predictions (similar to r-process studies). Especially a region of strong deformation and small $\\alpha$-capture Q-values around $A$=80 seems to be predicted with large variations among different mass models, and the even-even $N$=$Z$ nuclei between A=68 and 100 play a dominant role. \\nopagebreak" + }, + "9701/astro-ph9701092_arXiv.txt": { + "abstract": "We consider electron acceleration by obliquely propagating fast mode waves in magnetically dominated accretion disk coronae. For low coronal plasma densities, acceleration can exceed Coulomb drag at lower energies and energize electrons out of the thermal background, resulting in a nonthermal tail. The extent of this tail is determined by the balance between acceleration and radiative cooling via inverse Compton scattering and synchrotron emission, and usually goes out to tens of MeV. This will have direct applications for explaining the gamma-rays from several galactic black hole candidates, such as Cyg X-1 and GRO J0422, which show $0.5$--$5$ MeV emissions in excess over what most thermal models predict. Detailed time evolutions of the particle distributions and wave spectra are also presented. ", + "introduction": "Thermal Comptonization models have had much success in explaining the hard X-ray spectra from galactic black hole candidates (GBHCs) with the plasma temperature $\\sim 50$-$100$ keV and Thomson depths of a few (e.g., \\cite{sle76}; \\cite{st80}; \\cite{har94}; \\cite{lia93}). However, the most sensitive observations of GBHCs to date in the $0.5$--$5$ MeV range by {\\it Compton} Gamma-Ray Observatory have clearly revealed that persistent gamma rays ($>1$ MeV) are being produced in some GBHCs, notably Cyg X-1 (\\cite{mcc96}; \\cite{ph96}; \\cite{ling96}) and GRO J0422 (\\cite{dij95}). These gamma-ray emissions are very difficult to accommodate by the pure thermal models (e.g., \\cite{st80}; \\cite{t94}), strongly suggesting the need for modification (\\cite{sd95}) or to incorporate some nonthermal processes. Recently, Li, Kusunose, \\& Liang (1996) have proposed a gyroresonant stochastic electron acceleration model to account for the MeV emissions from GBHCs (see also Dermer, Miller, \\& Li 1996). In that model, they showed that high frequency whistlers can accelerate electrons directly from the thermal background, after which Alfv\\'en waves would continue the acceleration to higher energies. The result is an electron distribution with a hard non-Maxwellian tail. Compton scatterings from both thermal and nonthermal electrons produce a broad band X- to gamma-ray spectrum, in agreement with that observed from both Cyg X-1 and GRO J0422. One uncertainty associated with that model is the generation of whistlers, which can arise from either a cascade of wave energy from lower frequencies or a microinstability (see e.g., \\cite{gary93}). While the nature of wave cascading is fairly well known in the MHD regime (e.g., \\cite{v94}; \\cite{rw95}), it has not been investigated at higher frequencies. Producing waves that can gyroresonate with an electron of a given energy through a resonant microinstability requires (among other things) an anisotropic electron distribution containing electrons of that energy. Nonresonant instabilities may also produce the needed waves, but still require specific anisotropic distributions. In this {\\it Letter}, we consider electron acceleration by MHD fast mode waves that likely exist in accretion disk coronae. These low-frequency waves can accelerate electrons from thermal to relativistic energies. Coupled time-dependent diffusion equations for the electron distribution and the wave spectral density are solved numerically (\\cite{mlm96}, hereafter MLM). Nonthermal electron distributions are clearly obtained under a range of parameters in the accretion disk environment, which will have a direct bearing on the future modeling of the hard X-ray and gamma-ray emissions from GBHCs. ", + "conclusions": "We have studied particle acceleration in galactic black hole accretion disk coronae via interactions between electrons and fast mode waves---specifically, via the transit-time damping process. Including Coulomb collisions, inverse Compton scattering and synchrotron losses, we show that particles with speeds higher than the Alfv\\'en speed can be accelerated out of the thermal background, and we obtain steady state particle distributions composed of a Maxwellian plus a nonthermal high energy tail extending into several tens of MeV. Detailed radiation modeling will be presented in a forthcoming work and we expect that the Maxwellian and the nonthermal tail of the particle distribution will be responsible for the power-law spectra in the tens of keV and the high energy gamma-rays observed from several GBHCs such as Cyg X-1 and GRO J0422, respectively. The generation of plasma wave turbulence in accretion disk environments is a fairly unexplored topic, but it is reasonable to suppose that the fast mode waves will be excited since it is an intrinsic long-wavelength mode of a magnetized plasma. We emphasize that the coronal plasma $\\beta$ must be $< 1$ for electrons to get most of the wave energy, otherwise proton acceleration becomes possible, reducing the energy flow to the electrons. The particle acceleration mechanism discussed here also has direct implications on the high energy radiation from accretion disks in AGNs, notably Seyfert galaxies. Preliminary analyses have indicated that most of our results are insensitive to the size of the system; thus, we expect $>$ MeV emissions are also being produced in Seyferts as well, though high quality spectra above 200 keV are clearly needed to firmly settle this issue." + }, + "9701/astro-ph9701236_arXiv.txt": { + "abstract": "s{ STACEE is a proposed telescope for ground-based gamma-ray astrophysics between 25 and 500 GeV. The telescope will make use of large mirrors available at a solar research facility to achieve an energy threshold lower than existing ground-based instruments. This paper describes recent development work on STACEE.} ", + "introduction": "Discoveries from the Compton Gamma Ray Observatory (CGRO) \\cite{EGRET} and from ground-based experiments \\cite{Whipple} indicate that the high energy sky is rich with interesting astrophysics. Yet, there is a gap in experimental coverage between 20 and 250 GeV. Satellite instruments, such as GLAST,\\cite{GLAST} may eventually extend their reach above 20 GeV, but the experiments with the most promise to explore the gap in the near future are ground-based detectors using the atmospheric Cherenkov technique. The energy threshold of atmospheric Cherenkov detectors is governed by a number of parameters, of which the easiest to control is the mirror collection area.\\cite{Weekes} Large collection area translates into lower energy threshold, and large solar mirrors (heliostats) are readily available at existing power facilities. Since early 1994, we have been developing an experiment (STACEE) to use heliostat mirrors for Cherenkov astronomy. A similar experiment (CELESTE) is also under development in France.\\cite{Celeste} ", + "conclusions": "We have completed the design of an innovative atmospheric Cherenkov detector sensitive to gamma-rays in an unexplored energy region. The complete experiment can be built on a two year timescale." + }, + "9701/astro-ph9701222_arXiv.txt": { + "abstract": "In this article I review the main methods for determining the age of the Universe. I describe how to determine the age of the oldest known systems at $z=0$, the system of galactic globular clusters, using different techniques. I also describe how to date the Universe using the decay of radioactive elements (Cosmochronology). Finally, I focus on how to determine the age of the Universe at different redshifts and specially the age of radio-quiet galaxies at high redshift. I finish by arguing that the {\\it most probable age for the Universe is $14 \\pm 2$ Gyr}. ", + "introduction": "One of the most important questions in Cosmology is to know the age of the Universe. Unfortunately, a definitive answer is still to be found. Once the values of $\\Omega_0$ and $H_{0}$ have been established for our preferred cosmological model, the age of the Universe is unequivocally determined. In the case of an Einstein-de Sitter Universe $t=(2/3)H_{0}^{-1}$ (see e.g. Padmanabham 1993 for cases where $\\Omega \\ne 1$ and $\\Lambda \\ne 0$). On the contrary, if we are able to estimate the age of the Universe by other methods (e.g. dating stellar systems) then we can constrain the cosmological model and determine which kind of Universe we are living in. We can repeat the above argument at different redshifts, in this way we will gain information not only about $H_0$ but also about $\\Omega_0$ since, e.g., at $z=1.5$ the age of the Universe is $t=1.6h^{-1}$ Gyr if $\\Omega_0=1$ and $t=2.6h^{-1}$ Gyr if $\\Omega_0=0.2$, where $h=H/100\\, \\rm km s^{-1} Mpc^{-1}$. Therefore, an {\\it accurate} knowledge of the age of the Universe at different redshifts can tell us what kind of Universe we live in. Stars are the only clocks in the Universe available to us that are ``cosmology'' independent. The evolution of stars depends only on the rate at which the different nuclear burnings take place. So a good strategy to measure the age of the Universe is to find the oldest stars at every redshift. This article is organised as follows: In section \\ref{raul:stellar}, I review the stellar evolution of low mass stars. A detailed review on how to estimate globular clusters (GC) ages follows in the next section. Cosmochronology is the subject of the next section. In section \\ref{raul:galaxies}, I describe how to date high-redshift galaxies. I finish with a summary of the different ways to date the Universe. ", + "conclusions": "In this article we have focused our attention on how to calculate the age of the Universe at different redshifts and its cosmological implications. The main conclusions are the following: \\begin{itemize} \\item There are well defined stellar clocks in the Universe at different redshifts that allow us to obtain independent measures of the age of the Universe. \\item Galactic Globular Clusters are the best cosmological clocks at $z=0$. Several methods have been used to compute their ages. The {\\it Luminosity Function} is the most accurate method to compute the age and distance of Galactic Globular Clusters. {\\it The age of the oldest globular clusters is $14 \\pm 2$ Gyr}. \\item Cosmochronology gives a most probable age for the oldest stars of the Galaxy between 11 and 14 Gyr. \\item At high redshift quiet radio-galaxies are the best cosmological clocks to probe the early evolution of the Universe. I have described the case of 53W091 and give detailed description on how to compute its age. This galaxy (the oldest at $z=1.55$) was found to be 3.5 Gyr old. It represents a very rare event in modified cold dark matter models. \\end{itemize} Based on the previous arguments, we can say that the Age of the Universe is between {\\it 11 and 16 Gyr, with a most probable age of 14 Gyr}. This constrains the allowed values for $\\Omega$. If $H_0=65$ km s$^{-1}$ Mpc$^{-1}$ (see this volume) then $\\Omega \\leq 0.2$, otherwise $\\Lambda \\ne 0$. In the case $\\Omega=1$ then $H_0 \\approx 0.5$ km s$^{-1}$ Mpc$^{-1}$." + }, + "9701/astro-ph9701152_arXiv.txt": { + "abstract": "\\noindent In this paper we present preliminary results from our HST project aimed at exploring the connection between stellar dynamics and stellar evolution in the cores of high-density globular clusters. ", + "introduction": "Until a few years ago Galactic globular clusters (GC) were regarded as ideal laboratories for testing stellar evolution theories and for studying stellar dynamics in {\\it simple} stellar systems, and as important relics of the formation of the Galaxy. For many years most of the research proceeded as if the problems in the single fields mentioned above could be understood and solved independently of each other. Now, however, each branch of GC studies is realizing that further progress depends on viewing each cluster as a kind of {\\it ecosystem} of interrelated species. We have a growing body of observational evidence that dynamical interactions among stars in high-density clusters can modify their stellar content. Color gradients have been observed from the ground in the central regions of some post-core-collapse (PCC) or high-concentration clusters (\\cf\\ Djorgovski \\& Piotto 1993 for a review and references). The gradients are always in the sense of a bluer center, and extend even to the far-UV wavelengths. The gradients reflect radial changes in the stellar population. In addition, Fusi Pecci, Ferraro, \\& Cacciari (1993) have shown that the length of the horizontal branch (HB), and the presence and the extent of blue tails in particular, are correlated with the cluster density and concentration, in the sense of more concentrated or denser clusters having bluer and longer HB morphologies. The theoretical understanding of these phenomena remains unclear, though it is very likely that binaries and stellar interactions are involved in modifying stars located on the evolved branches of the color--magnitude diagram (CMD). The exceptional resolving power of HST is of fundamental importance in this kind of study, as it allows observing faint stars in the center, down to the main sequence. As an example, in Figure~1 we compare one of the best ground-based images on which our previous investigation (Djorgovski \\& Piotto 1993) was based with the corresponding HST frame. \\medskip \\hbox{ \\vbox{\\hsize 2.5 truein \\psfig{figure=piotto1.ps,height=2.6truein,width=2.6truein} } \\vbox{\\hsize 2.5 truein \\psfig{figure=piotto2.ps,height=2.6truein,width=2.6truein} } } \\cptn{Figure 1. The {\\it left panel} shows a ground-based {\\it V\\/}-band image of the inner $21\\times21$ arcsec$^2$ of the GC M30, taken at the ESO NTT telescope (seeing FWHM = 0.6 arcsec). The {\\it right panel} shows the same field from the PC frame of the WFPC2 camera in the F555W band.} \\bigskip In this paper we present some preliminary results from our HST project specifically devised to explore the connection between stellar dynamics and stellar evolution by investigating the GC star population in the very inner regions of 10 high-density clusters. ", + "conclusions": "" + }, + "9701/astro-ph9701042_arXiv.txt": { + "abstract": "Star clusters with a high central density contain an ecological network of evolving binaries, affected by interactions with passing stars, while in turn affecting the energy budget of the cluster as a whole by giving off binding energy. This is the first paper in a series aimed at providing the tools for increasingly realistic simulations of these ecological networks. Here we model the core of a globular star cluster. The two main approximations are: a density of stars constant in space and time, and a purely single star population in which collisions between the evolving stars are modeled. In future papers in this series, we will relax these crude approximations. Here, however, they serve to set the stage before proceeding to the additional complexity of binary star interactions, in paper II, and background dynamical evolution, in later papers. ", + "introduction": "In dense stellar systems, such as open and globular clusters and galactic nuclei, encounters between individual stars and binaries can affect the dynamical evolution of the system as a whole on a time scale comparable to, or even shorter than, a Hubble time. In order to reach a detailed theoretical understanding of such systems, the following three steps are necessary. First, we need to understand the basic mechanism of the dynamical evolution, in the limit of a point-mass approximation for the stars. Second, effects of dynamical encounters on the internal evolution of single stars and binaries has to be taken into account. Third, we have to model the feedback of these internal changes onto the dynamical evolution of the whole system. Let us briefly review each step. Great progress has been made with the first step, modeling the dynamical evolution of point-mass systems. In the seventies, the processes of core collapse and mass segregation were studied with the use of various types of Fokker-Planck approximations. In the eighties, these simulations were extended successfully beyond core collapse, and various studies were made of the phenomenon of gravothermal oscillations, ubiquitous in the post-collapse phase. Of these models a few even include mass loss due to the evolution of the stars (\\cite{cw90}). In the nineties, we are finally beginning to switch over from Fokker-Planck approximations to much more detailed and realistic $N$-body simulations. In 1995, the construction of the GRAPE-4, a special-purpose machine with Teraflops speed, has made a $32,000$--body simulation feasible, providing the first direct evidence of gravothermal oscillations (Makino, 1996a,b). Extending these simulations to the full realm of globular clusters ($N=10^5\\sim10^6$) will require Petaflops speed, something that could be realized by future special-purpose machines in the GRAPE series by as early as the year 2000. While the point-mass approximation provides a good qualitative guide for the construction of dynamical models of dense stellar systems, this approximation quickly breaks down when we require quantitatively accurate results. The second step attempts to model the effects of close encounters. A number of different investigations have estimated the rate at which physical collisions have take place, under various circumstances (\\cite{hd76}, \\cite{vm88}, \\cite{dr92} and \\cite{db95}). However, little progress has been made so far in following the changes induced in the stellar population, beyond enumerating the number of mergers. In the simulations presented below, collisions are modeled in an evolving population of single stars in a high-density stellar environment. Paper II in this series will extent our treatment to follow the induced changes in binary systems, both on the level of changes in orbital parameters as well as in the internal structure of the stars. However, such investigations are only a start, and cannot lead to a quantitative modeling of dense stellar systems, since they are not yet self-consistent. What is needed in addition is a treatment of the feedback mechanism, from the changes in single stars and binaries back to the overall dynamics of the system. This third step is being pioneered for open clusters by \\cite*{aar96}. The current series of papers aims to provide self-consistent models of this type, by coupling relatively crude stellar evolution recipes, documented and tested in papers I and II, to a fully dynamical $N$-body system. This paper is organized as follows. Our approach to the study of the ecology of star clusters is summarized in somewhat more detail in \\S 2. The next section, \\S 3, describes our simulation techniques, and the various approximations involved. In \\S4, we present the result of a simulation starting a single model, with a minimum of free parameters. The results of a more realistic core model run are presented in \\S5. \\S6 sums up. ", + "conclusions": "The models discussed in this paper are very crude in their treatment of the encounter processes, of the result of a collision between two stars, and of the evolution of the merger products. Apart from these approximations and the fact that we use a stellar evolution model for population~I instead of pop.~II stars, the adopted mass function is also highly uncertain. Nonetheless, some interesting results can be delineated. Comparison with the calculations of \\cite*{db95} shows the effect of allowing the merger products to evolve. An immediate consequence of this is the lower prediction for the number of blue and yellow stragglers present in the cluster (as is clear from Figure~\\ref{fig_nbgss}). The formation rates of blue and yellow stragglers give a poor indication for the actual number of stragglers present in the cluster at a particular instant. Due to the low density of model $S$ the collision frequency is small. The steep Salpeter mass function also suppresses the encounter rate and the production of stellar curiosities; the majority of the collisions involve two rather low mass main-sequence stars which results in a merger that evolves too slow to produce a blue stragglers within the time span of the simulation. The Hertzsprung-Russell diagram of our model cluster (model $C$) shows that blue stragglers close to the turn-off point lie on the main sequence, whereas blue stragglers above the turnoff point are mostly found at some distance from the main sequence. The reason for this is that collisions only become important in the cluster when an initial period of low density is followed by the contraction of the cluster core. The more massive blue stragglers are formed in collisions between stars close to the terminal-age main sequence, and evolve relatively quickly. Blue stragglers close to the turn-off are formed in collisions between relatively low-mass stars which did not evolve very far away from the zero-age main sequence, and therefore also the merger products are close to the zero-age main sequence, and evolve slowly. Thus, the point where blue stragglers have left the main sequence gives an indication of the time when collisions in the cluster became frequent (see also \\cite{pz96a}). Our model $C$ predicts a depletion of giants, in the core only, up to $\\sim 50\\%$ shortly after $t_{\\rm cc}$ relative to a collision-less stellar system, in globular clusters with a collapsed core where the fraction of horizontal-branch stars is enhanced. Consequently the depletion of giants relative to the number of horizontal branch stars is strongly present in the high-density stellar system. Collisions between single stars cannot explain the observation that giants can be depleted well outside the core or completely absent in it, as observed in the core of M~15 (\\cite{dpc+91}). In our simulated cluster cores the total number of white dwarfs that exceed the Chandrasekhar limit due to accretion from a circum-stellar disc is small, even in the cluster simulation with the highest density. In model $C$ 8\\%\\ of the white dwarfs experience an accretion-induced collapse, which (after correction for the ratio $f_{\\rm c}$ between the number of stars in the model and in an actual core -- see Table~\\ref{Tab_init}) corresponds to 190 supernovae of type Ia during the 6$\\,$Gyr of our calculation. If all of these collapses would lead to the formation of a neutron star, and if all of these would remain in the core, this would be a substantial addition to the total number of neutron stars in the core, which is about 460 (after correction for $f_{\\rm c}$) at the start of our calculation. This result, however, strongly depends on the adopted mass function for the white dwarfs. The formation-rate of neutron stars with an accretion disc and the subsequent formation of a recycled pulsar or black hole is (to first order) linearly dependent on the number of neutron stars, which depends not only on the initial mass function but also on the subsequent mass segregation in the cluster. The encounter rates between neutron stars and main-sequence stars are similar in our calculations to the rates found in the calculations by \\cite*{vm88}, by Di~Stefano \\& Rappaport (1992) and by Davies \\&\\ Benz (1995). \\nocite{vm88}\\nocite{db95}\\nocite{dr92} After a collision between a neutron star and another cluster member the merged object becomes visible as an X-ray source (for at most 1~Gyr) after which it becomes a recycled pulsar or, if its mass exceeds 2~\\msun, a black hole. The total number of such X-ray sources, recycled pulsars or black holes scales linearly, in first order, with the number of neutron stars in the cluster core, which is rather uncertain. Our computations reveal that collisions between single stars result in a small number of recycled pulsars: about 70 are formed in a core (after correction for $f_{\\rm c}$) according to model $C$. Whether this is enough to explain the observed numbers is not clear. The intrinsic luminosity distribution of millisecond pulsars, and hence the fraction of them that is detectable in a typical cluster, is not known; and some clusters with high encounter rates show remarkably few recycled pulsars, the globular cluster NGC~6342 is an example (see \\cite{lyn93})." + }, + "9701/astro-ph9701104_arXiv.txt": { + "abstract": "The most significant periodicities in the terrestrial impact crater record are due to the \\signal: the bias of assigning integer values for the crater ages. This bias seems to have eluded the proponents and opponents of real periodicity in the occurrence of these events, as well as the theorists searching for an extraterrestrial explanation for such periodicity. The \\signal ~should be seriously considered by scientists in astronomy, geology and paleontology when searching for a connection between terrestrial major comet or asteroid impacts and mass extinctions of species. ", + "introduction": "An outstanding series of papers appeared in 1984 when a 28.4~Myr cycle was detected in the terrestrial impact crater record (Alvarez \\& Muller 1984, Davis et al. 1984, Whitmire \\& Jackson 1984). This value was close to the 26~Myr cycle discovered in the geological record of major mass extinctions of species (Raup \\& Sepkoski 1984). The fascinating idea of periodic comet impacts causing ecological catastrophies emerged (Alvarez \\& Muller 1984, Davis et al. 1984). It was suggested that an unseen solar companion (Nemesis) might induce gravitational disturbances to the Oort comet cloud triggering periodic cometary showers (Davis et al. 1984, Whitmire \\& Jackson 1984). Other astronomical models have been proposed later to account for the above periodicities, the ``galactic carrousel'' being perhaps the most widely accepted model (e.g. the review by Rampino \\& Haggerty 1996, and references). The main idea of the ``galactic carrousel'' model is that the Oort comet cloud is periodically perturbed by galactic tides as the Solar System revolves around the centre of the Milky Way galaxy. We will show that only one extremely significant regularity exists in the impact crater record: the \\signal. ", + "conclusions": "A few topics must be emphasized to avoid misunderstanding. ({\\sc{i}}) We did not assume that the integer $t_{\\rm i}$ cause the detected periodicities. On the contrary, the periodicities were uniquely detected with non--parametric statistical methods (i.e. model independent). {\\it We simply tested the ``null hypothesis'' ($H_0$) that the impact crater ages represent a random sample of circular data.} The \\signal ~reaches $\\gamma \\!=\\!0.001$ in all subsamples \\crione, ..., \\criseven, the critical levels of the \\SHORTSD ~and {\\WSD}s being extremely high. The analytical statistics of our methods are robust, i.e. Monte Carlo or other computational techniques are unnecessary. ({\\sc{ii}}) The \\signal ~induces irregular multi--modal $\\phi_{\\rm i}$ distributions, which were not detected earlier with methods sensitive to uni--modal $\\phi_{\\rm i}$ distributions. For example, the power spectrum method is most sensitive to sinusoidal variations. If the $\\phi_{\\rm i}$ distribution for the ``correct'' $P'$ were exactly bi--modal, the peaks of the power spectrum would be at $P'$, $P'/2$, ... But the power spectrum method is quite insensitive to more irregular $\\phi_{\\rm i}$ distributions. That the \\signal ~was not detected earlier, over twelve years after the study by Alvarez \\& Muller (1984), is a direct consequence of favouring methods sensitive to uni--modal $\\phi_{\\rm i}$ distributions. Why should the $\\phi_{\\rm i}$ distributions connected to the possible periodicity in terrestrial impact cratering rate or mass extinctions of species actually be uni--modal or of any regular shape? ({\\sc{iii}}) It may well be that the large $\\sigma_{\\rm t_i}$ prevent detection of periodicity (e.g. Heisler \\& Tremaine 1989). In that case our study is simply an exercise of statistics providing one new argument against real periodicity. However, considering the prevailing theories based on assuming the presence of real periodicities (see e.g. Rampino \\& Haggerty 1996: ``Shiva Hypothesis''), it was a high time to perform this exercise. ({\\sc{iv}}) The \\signal ~has most probably induced those spurious periods with more or less unimodal $\\phi_{\\rm i}$ distributions, which were detected in several earlier studies (e.g. Alvarez \\& Muller 1984, Yabushita 1991). In any case, the \\signal ~is clearly the most significant periodicity in the impact crater record. ({\\sc{v}}) We do not argue that the major comet or asteroid impacts and the mass extinctions of species are uncorrelated, but emphasize that the \\signal ~dominates the time distribution of the former events. Our conclusions are simple. The epochs of mass extinction events of species may follow a possibly ``nonhuman'' cycle of 2.76~Myr, but the currently available impact crater data definitely reveals the embarrassing \\signal. The fellow scientists have unconsciously offered a helping hand to the Nemesis (e.g. Davis et al. 1984) or ``galactic carrousel'' (e.g. Rampino \\& Haggerty 1996). The arduous task for the future geological research is to determine more accurate (preferably noninteger) revised ages for impact craters to eliminate the \\signal, which may then lead to a detection of real periodicity. Over a decade has elapsed in redetecting the regularities of our own integer number system and then interpreting them as periodicity in the ages of impact craters. \\vspace{-0.2cm}" + }, + "9701/astro-ph9701038_arXiv.txt": { + "abstract": "We have obtained near-infrared observations of some of the faintest objects so far known towards the Pleiades young stellar cluster, with the purpose of investigating the sequence that connects cluster very low-mass stars with substellar objects. We find that infrared data combined with optical magnitudes are a useful tool to discriminate cluster members from foreground and background late-type field stars contaminating optical surveys. The bottom of the Pleiades sequence is clearly defined by the faint HHJ~objects as the very low-mass stars approaching the substellar limit, by the transition object PPl~15, which will barely ignite its hydrogen content, and by the two brown dwarfs Calar~3 and Teide~1. Binarity amongst cluster members could account for the large dispersion observed in the faint end of the infrared colour-magnitude diagrams. Two objects in our sample, namely HHJ~6 and PPl~15, are overluminous compared to other members, suggesting a probable binary nature. We have reproduced the photometric measurements of both of them by combining the magnitudes of cluster very low-mass stars and brown dwarfs and using the most recent theoretical evolutionary tracks. The likely masses of the components are slightly above the substellar limit for HHJ~6, while they are 0.080 and 0.045$\\pm$0.010~$M_{\\odot}$ for PPl~15. These masses are consistent with the constraints imposed by the published lithium observations of these Pleiads. We find a single object infrared sequence in the Pleiades cluster connecting very low-mass stars and brown dwarfs. We propose that the substellar mass limit ($\\sim$0.075~$M_{\\odot}$) in the Pleiades ($\\sim$120~Myr) takes place at absolute magnitudes $M_{\\rm I}$~=~12.4, $M_{\\rm J}$~=~10.1, $M_{\\rm H}$~=~9.4 and $M_{\\rm K}$~=~9.0 (spectral type M7). Cluster members fainter by 0.2~mag in the $I$-band and 0.1~mag in the $K$-band should be proper brown dwarfs. The star-brown dwarf frontier in the Hyades cluster (600~Myr) would be located at $M_{\\rm I}$~=~15.0, $M_{\\rm J}$~=~11.6, $M_{\\rm H}$~=~10.8 and $M_{\\rm K}$~=~10.4 (spectral type around M9). Fon an age older than 1000~Myr we estimate that brown dwarfs are fainter than $M_{\\rm K}$~=~10.9 (spectral type later than M9.5). \\keywords {Stars: pre-main sequence -- Stars: late-type -- Stars: low-mass, brown-dwarfs -- Open clusters: Pleiades} ", + "introduction": "The $K$ vs ($I-K$) colour-magnitude diagram for our sample (filled symbols) and some of the Pleiades least massive objects discovered to date (open symbols) is shown in Fig.~1. We also plot previously known proper motion cluster members observed at these wavelengths by Steele et al. (\\cite{steele93}) using the revised $I$ photometry published in Steele \\& Jameson (\\cite{steele95}). We have also included the BD candidates proposed by Stauffer et al. (\\cite{stauffer89}) and Williams et al. (\\cite{williams96}). Different symbols are used for clarity and no de-reddening has been applied. The solid line is an averaged MS derived from the photometry of field stars taken from Leggett (\\cite{leggett92}), shifted to an assumed Pleiades distance modulus of ($m-M$)$_{\\circ}$~=~5.53~mag and adopting $A_{\\rm V}$~=~0.12~mag (Crawford \\& Perry \\cite{crawford76}); $A_{\\rm I}$ and $A_{\\rm K}$ were then derived using the relationships of Rieke \\& Lefobski (\\cite{rieke85}). \\begin{figure} \\psfig{figure=oso_f1.ps,width=8cm,angle=0.0} \\caption[]{$K$ vs ($I-K$) colour-diagram for our sample and other Pleides low-mass objects. Filled circles denote cluster members whereas filled squares denote non-members (see text for explanation). Asterisks stand for the BD candidates of Williams et al. (1996), and triangles for those of Stauffer et al. (1989). Proper motion HHJ objects (data taken from Steele et al. 1993, and Steele \\& Jameson 1995) are shown as open circles. The line represents the MS as derived from the photometry of field late-type dwarf stars (Leggett 1992), shifted to the Pleiades distance and reddening. Error bars for our photometry are indicated at the top right corner} \\end{figure} Only objects with spectral types later than M3 (($I-K$) $\\ge$ 2.2) are displayed in Fig.~1. As expected, young Pleiades members have not reached the field MS (Stauffer \\cite{stauffer84}; Steele \\& Jameson \\cite{steele95}), thus lying above it. This is basically what is observed except for a few exceptions, like PPl~3 (Stauffer et al. \\cite{stauffer89}), also named as JS1 (Jameson \\& Skillen \\cite{jameson89}), which is located 1.7~mag below the MS line. The infrared photometry rules out its membership in the cluster, a fact that is consistent with our proper motion measurement in Paper~I. Five very low-mass stars proposed by Williams et al. (\\cite{williams96}) around ($I-K$)~=~2.5 and $K\\sim14$ lie also below the field MS and are indeed fainter in the $K$-band than proper motion Pleiades of the same colour, suggesting that they are unlikely members based on their infrared photometry. Two other objects in our sample, Calar~2 and Calar~5, are clearly located below the MS by $\\sim$1~mag in Fig.~1. Both of them have infrared colours that resemble those of field stars of similar spectral types within the errors in photometry and classification (e.g., see the ($I-K$) vs spectral type diagram of Fig.~2). Although found with ($R-I$) values and radial velocities consistent with other cluster members, Calar~2 and Calar~5 are non-members on the basis of their infrared photometry, confirming the suspicions of Mart\\'\\i n et al. (\\cite{martin96a}). These objects do not fit in the sequence described by the cluster members in Fig.~4 in Mart\\'\\i n et al.'s work, being fainter in $I$ than expected for Pleiads of these spectral types. Calar~2 and Calar~5 are therefore likely background reddened low-mass stars. Infrared observations allow an easy discrimination of contaminants that could arise in optical photometric surveys from true Pleiads. Roque~1 appears very close to the MS in Fig.~1, but slightly below it implying that it is a likely non-member. By inspecting Fig.~2 and allowing for the photometric uncertainties, this object fits in the ($I-K$) colour-spectral type diagram for solar metallicity field stars. From the comparison of its photometry and spectral type with the absolute magnitudes published in Kirkpatrick \\& McCarthy (\\cite{kirkpatrick94}) we conclude that this object may be located at approximately the Pleiades distance. Nevertheless, given its high radial velocity ($v_{\\rm r}$~=~151$\\pm$30~km s$^{-1}$), it was suggested in Mart\\'\\i n et al. (\\cite{martin96a}) that Roque~1 could be part of the old galactic population, but neither its $R$, $I$ and $J$, $H$, $K$ colours nor its optical spectrum show any evidence of significant low-metallicity. Roque~1 still awaits further investigation that confirms its membership in an old population. Calar~1 sits on the MS in Fig.~1. However, its location is highly unexpected given its spectral type and ($R-I$) colour. This object was found to be the reddest BD candidate in the optical survey of Paper~I with ($R-I$)~=~2.9, and surprisingly, it appears rather blue in the infrared (see Fig.~2). Apart from being one of the few M9 dwarfs known in the whole sky, it is the first one to our knowledge that behaves so rarely at infrared wavelengths. Leggett (\\cite{leggett92}) states that for very cool dwarfs no obvious metallicity effects are discernable in optical diagrams, but that they become apparent in infrared plots. Metal-poor objects are well known to have bluer colours than solar metallicity ones. Calar~1 may be a metal-deficient dwarf located foreground to the Pleiades, a fact that is supported by its high radial velocity ($v_{\\rm r}$~=~85$\\pm$30~km s$^{-1}$, Mart\\'\\i n et al. \\cite{martin96a}). The proper motion and parallax determinations would help to understand the nature of this object. Calar~1, Calar~2, Calar~5 and Roque~1 are excluded from further discussion in this paper as very unlikely members of the cluster. \\begin{figure} \\psfig{figure=oso_f2.ps,width=8.8cm,angle=-90} \\caption[]{($I-K$) vs spectral type diagram for our sample. Symbols are as in Fig.~1. Spectral types have been taken from Steele \\& Jameson (1995), Mart\\'\\i n et al. (1994), and Mart\\'\\i n et al. (1996). The line represents the relationship for solar metallicity field stars (Kirkpatrick \\& McCarthy 1994). Error bars are half a spectral subclass (as stated by the authors), and $\\pm$0.1~mag for the photometric colour. Some objects have been labeled for clarity} \\end{figure} The remaining objects in our sample (HHJ~2, HHJ~3, HHJ~6, PPl~1, PPl~15, Teide~1 and Calar~3) do follow the sequence described by the proper motion cluster members in Fig.~1. Particularly, Teide~1 and Calar~3 define the present faint end of the Pleiades, which lies 0.3--0.45~mag above the MS in the $K$-band. Objects with similar photometry should be considered as probable BDs in the cluster. The object in Williams et al. (\\cite{williams96}) labelled as No.~13 could be a BD based on its photometry. However, we have recently obtained deep images in the $I$-band which show that it is likely an extense object. \\subsection{Broad-band energy distributions} We have combined optical ($R$ and $I$) and infrared magnitudes of the Pleiades objects in our sample in order to obtain the general shape of their energy distributions and compare them with those of field stars with similar spectral types. We adopted the absolute magnitudes for M dwarfs given in Kirkpatrick \\& McCarthy (\\cite{kirkpatrick94}) as representative of the field, and the distance modulus and extinction of the Pleiades to convert the observed magnitudes to absolute magnitudes for the cluster members. The zero-magnitude fluxes and center-wavelengths (Cousins for $R$ and $I$, UKIRT for $J$, $H$ and $K$) are taken from Mead et al. (\\cite{mead90}). In Fig.~3 we plot the final broad-band spectral distributions for PPl~1, PPl~15, Teide~1, and Calar~3. PPl~1 is the only object in our sample which has no $R$ magnitude available. We have estimated it using the mean ($R-I$) colour of the other three M6.5 Pleiads in our sample (the mean coincides with that given by Kirkpatrick \\& McCarthy \\cite{kirkpatrick94} for a field star of the same spectral type). Overplotted are the distributions of stars from M5.5 to M9 (for spectral types M8 and M9 Kirkpatrick \\& McCarthy \\cite{kirkpatrick94} do not provide data at the $R$-band). This plot clearly shows how important infrared emission is to study these extremely late-type objects since most of their bolometric flux is emitted at these wavelengths. Less than a few per cent is emitted blueward of the $R$-band. Given a spectral type, a Pleiades member having the same effective temperature and hence a very similar energy distribution than a field star, should emit more at each wavelength due to its youth; i.e., it should move up along the flux-axis in Fig.~3. PPl~1 and PPl~15 appear to have energy distributions that resemble within uncertainties in flux conversion and observed magnitudes those of the old M6--M7 stars, except that the formers (which have similar optical spectral types) are twice (0.3~dex) more luminous. Teide~1 and Calar~3 show cooler energy distributions that are much alike to M8-type field stars, but the two Pleiades BDs are overluminous by a factor 1.6 (0.2~dex). These results agree with expectations for cluster members. \\begin{figure} \\psfig{figure=oso_f3.ps,width=8cm,angle=0.0} \\caption[]{Broad-band energy distributions at $R$, $I$, $J$, $H$ and $K$ wavelengths. Pleiades members are denoted as circles and full lines. The spectral distributions of M5.5--M9 field dwarfs are also plotted for comparison (crosses and dotted lines). Errors in flux conversion could be as large as 20\\%.} \\end{figure} The accurate determination of bolometric luminosities for our objects would require the use of spectroscopic data covering the whole optical and infrared regions. Integrating the broad-band flux distributions always overestimates the luminosity, because the spectra of these late-type objects are dominated by deep molecular and absorption features that are not precisely described by the photometry. The overestimation (less than 10\\% \\ according to Reid \\& Gilmore \\cite{reid84} and Tinney et al. (\\cite{tinney93}) seems to be constant through all the spectral types, although perhaps slightly larger for M9 stars. This allow us to be confident on the fact that relative luminosities within our sample are not strongly influenced by molecular absorption effects. We have integrated the flux distributions of our Pleiades members and evaluated the differences in luminosity with respect to Teide~1 (the least luminous object in our sample). The resulting values are provided in Table~2. We estimate that they are affected by an error of $\\pm$0.04~dex which mainly comes from uncertainties in the photometry. \\begin{table} \\caption[]{Differences in log~($L/L_{\\odot}$) relative to Teide~1} \\begin{center} \\begin{tabular}{cccccc} \\hline HHJ~6 & HHJ~2 & HHJ~3 & PPl~1 & PPl~15 & Calar~3 \\\\ \\hline 0.63 & 0.41 & 0.45 & 0.36 & 0.37 & 0.05 \\\\ \\hline \\end{tabular} \\end{center} \\end{table} In the Pleiades cluster stars of spectral type M6.5 are 2.29$\\pm$0.25~times (0.36$\\pm$0.04~dex) more luminous than M8-type objects. We have computed the mean luminosities for field stars of the same spectral types averaging the data from Tinney et al. (\\cite{tinney93}) and Bessell \\& Stringfellow (\\cite{bessell93}), and found that M6.5 field stars are overluminous by a factor 1.78 (0.25~dex) with respect to M8 stars. It is remarkable that the factor for the Pleiades is larger by $\\sim$0.1~dex. This offset in luminosity takes into account differences in radii. The fact that the ratio $R_{\\rm M6.5}/R_{\\rm M8}$ is greater in the Pleiades cluster than in the field could be explained by a larger mass difference between a M6.5 object and a M8 object in the former than in the latter. This result was expected from theoretical evolutionary models. It would be desirable to define a consistent spectral type classification using objects that belong to a cluster to take advantage of the coeval formation as well as of the homogeneous metallicity. ", + "conclusions": "We have presented near-infrared $J$, $H$ and $K$ photometry for some of the least luminous objects so far known in the Pleiades cluster. They are very low-mass stars approaching the substellar limit, transition objects, and massive BDs (HHJ and PPl objects, and Calar~3 and Teide~1). In addition, we have also included in our sample four of the BD candidates proposed in Paper~I (Calar~1, 2, 4, 5 and Roque~1) in order to asses their membership on the basis of the infrared data. We have made use of all the available spectroscopic and photometric knowledge on our final list of objects to characterize the sequence that cluster members around the substellar limit and beyond do follow in the infrared colour-magnitude diagrams. This sequence is clearly defined by the faintest HHJ~objects, PPl~1, PPl~15, Calar~3 and Teide~1, lying 0.3--0.4~mag above the field MS for spectral types later than M6. Objects found in future surveys with similar photometry (optical and infrared) should be considered as likely cluster members. Undoubtely, infrared data is highly useful to discriminate those objects that are true Pleiades members from those that are not. Optical surveys could be contaminated by background objects that although fitting in the cluster sequence at these wavelengths, may suffer from interestellar reddening. As extinction affects much less infrared magnitudes, consequently these objects do not match the $J$, $H$ and $K$ photometry expected for the Pleiades (this is the case of Calar~2 and Calar~5). Foreground late-type stars are also possible contaminants in photometric surveys, although its contamination is expected to be considerably smaller. However, in the optical survey of Paper~I two of these field stars were detected, Roque~1 and Calar~1, both having very high radial velocities (Mart\\'\\i n et al. \\cite{martin96a}). The reddest object in optical wavelengths, Calar~1 (M9), appears to have rather bluer infrared colours than expected for its spectral type. This together with its high radial velocity suggests membership to an old population. The large dispersion observed in the infrared photometry of Pleiades members is likely attributable to binarity effects. Two objects in our sample, namely HHJ~6 and PPl~15, lie $\\sim$0.7~mag above the single cluster sequence in infrared diagrams. Adopting an age of 120~Myr for the cluster and using the photometric data of very low-mass Pleiades stars and BDs as well as the theoretical tracks of Chabrier et al. (\\cite{chabrier96}), we have reproduced the photometry of both, finding that the likely masses of the components are 0.085 and 0.085$\\pm$0.010~$M_{\\odot}$ for HHJ~6, 0.080 and 0.045$\\pm$0.010~$M_{\\odot}$ for PPl~15. These results are in fairly good agreement with those obtained from the lithium test. Follow-up infrared spectroscopic observations of PPl~15 should be made to try to detect the effects of the companion on the spectrum, which are expected to be detectable. Since the primary of PPl~15 would be exactly located at the transition region between stars and BDs in the Pleiades, we have used it to define the empirical location of the substellar limit in infrared diagrams, for which we have corrected the observed magnitudes from the contribution of the secondary. We have estimated that the star-BD frontier on the single Pleiades sequence is located at magnitudes $I$~=~17.8--18.0, $J$~=~15.6--15.7, $H$~=~14.9--15.0 and $K$~=~14.5--14.6 and at optical spectral types M6.5--M7. New Pleiades objects that are found fainter along the photometric sequence defined in this paper might be BDs indeed. Based on this result, we estimate that the hydrogen-burning mass limit occurs at spectral types around M9 in the Hyades cluster." + }, + "9701/astro-ph9701110_arXiv.txt": { + "abstract": "Spheroidal components of spiral galaxies have been considered the only dynamically important component in gravitational lensing studies thus far. Here we point out that including the disk component can have a significant effect, depending on the disk inclination, on a variety of lensing properties that are relevant to present studies and future surveys. As an example, we look at the multiple image system B1600+434, recently identified as being lensed by a spiral galaxy. We find that including the disk component one can understand the fairly large image separation as being due to the inclination of a typical spiral, rather than the presence of a very massive halo. The fairly low magnification ratio can also be readily understood if the disk is included. We also discuss how such lensed systems might allow one to constrain parameters of spiral galaxies such as a disk--to--halo mass ratio, and disk mass scale length. Another example we consider is the quasar multiple--lensing cross section, which we find can increase many--fold at high inclination for a typical spiral. Finally, we discuss the changes in the gravitational lensing effects on damped Lyman alpha systems (DLAS) when disk lensing is included. ", + "introduction": "Galaxies responsible for gravitational lensing have typically been modeled as one-component systems in which the entire mass distribution of the galaxy lens, visible plus dark matter, is fairly well approximated by a singular isothermal sphere (SIS): $\\rho(r)=\\rho_o (r_o/r)^2$ (see e.g. Narayan and Bartelmann \\markcite{NB}1997). This is probably justified for E/SO galaxies where both components have a surface density fairly independent of orientation (a small ellipticity is required only to understand the presence of four-image systems) and the high central density represents the high density of the visible matter. By contrast, in spiral galaxies the disk component has a very different structure from the dark matter halo and the bulge, which results in significantly different projected surface density depending on orientation. Given that the projected surface density of the disk changes by a factor of radius/thickness $\\sim 15$ between a face-on and an edge-on spiral, and that the dark and visible components are known to contribute roughly equally to the total mass within 4 -- 5 disk scale lengths, one would expect significant effects on lensing properties for images within a couple of optical radii. The identification of the lensing galaxy in the double image system B1600+434 as a nearly edge on spiral by Jaunsen and Hjorth \\markcite{JH}(1997) gives us a perfect example to demonstrate the importance of including the disk component in modeling lensing by spirals. However, an understanding of the lensing geometry of spirals is important for several reasons besides that of modeling observed systems. First, recently the gravitational lensing effects on damped Lyman alpha systems (DLAS) have been investigated (Smette, Claeskens, and Surdej\\markcite{SCS} 1997; Bartelmann and Loeb \\markcite{BL}1996). At least some DLAS are believed to be due to absorption in protogalactic disks or spiral galaxies (Wolfe \\markcite{W}1988, \\markcite{W2}1997), and if the gas column density falls off exponentially, one is necessarily dealing with lensing within a few exponential scale lengths of the galaxy center, which is where we expect disk lensing to be important. Secondly, the extra parameters added by introducing a disk component are the physically interesting ones of disk surface density and scale length, so that the observational study of gravitational lensing by spirals will provide new constraints on the distribution of mass in the various components of spirals. Although only a few such systems are currently known, surveys such as the Sloan Digital Sky Survey will catalog a large number of quasars, among which there should be a significant number lensed by spiral galaxies. This is why it is also important to understand the effect of disk lensing on statistical properties such as the cross section for multiple-image lensing by spirals. In this paper we first explain our method of modeling spiral galaxy lenses as singular isothermal spheres plus infinitesimally thin disks (SIS+disk). We start with the simple example of a constant density disk, to show the effect of adding a disk. We show that such a model fits many of the observations on B1600+434, and discuss the results of modeling, more properly, with an exponential disk. We then discuss the dependence of the lens cross section for multiple lensing of background sources on the inclination of the disk, and discuss also how it would enhance the gravitational lensing effects on DLAS. We finish with the conclusions and implications from our work. ", + "conclusions": "We have shown that the inclusion of a disk component in a spiral gravitational lens model has significant effects that cannot be ignored. Our results imply that a host of issues ranging from the ability of spirals to lens QSO's to detailed properties of lens systems should be reconsidered to properly include the lensing effects of their disks. We have shown, as an example, that a SIS+disk model can account for the observations of the B1600+434 system with fairly reasonable parameters. The implications of our results are that a velocity dispersion $\\sigma_v$ significantly smaller than the SIS prediction of 200 km/s, and a time delay greater than one month, are to be expected. With more data and more detailed analysis, the modeling of multiply--imaged systems lensed by spiral galaxies will constrain the relative amounts of matter in the disk and halo when disk lensing is included. However, because of the dependence on new parameters, further study will be necessary before such systems can give an unambiguous measurement of the Hubble parameter $h$. We have also noted that the disk contribution to lensing will make the amplification and by-pass effects on the observed distribution of DLAS in bright QSO's even more pronounced at high column densities. After the first version of this paper was posted on the preprint archive as astro-ph/9701110, the preprint astro-ph/9702078 was posted by Wang and Turner discussing critical curves and caustics of inclined, constant density disks. Their equations and figures can be seen to be the same as ours if one is careful to scale appropriately by the Hubble constant $h$, and to notice that they have calculated angular diameter distances with all the matter in the universe in clumps while we have assumed it to be distributed smoothly. This work has been supported by NSF grant PHY-9600239 at UM-St. Louis, and by NSF grant PHY-9402455 and a NASA ATP grant at UCSC. AM also acknowledges GAANN support. The authors are grateful for helpful suggestions from the referee, Matthias Bartelmann; for help from Raja Guhathakurta and Steve Vogt in interpreting the HST archival image of B1600+434; and correspondence from A.O. Jaunsen, Chris Kochanek, and Tom Schramm. \\clearpage" + }, + "9701/astro-ph9701056_arXiv.txt": { + "abstract": "The energy loss rate of a magnetized electron gas emitting axions $a$ due to the process $e^- \\to e^- +a$ is derived for arbitrary magnetic field strength $B$. Requiring that for a strongly magnetized neutron star the axion luminosity is smaller than the neutrino luminosity we obtain the bound $g_{ae} \\lesssim 10^{-10}$ for the axion electron coupling constant. This limit is considerably weaker than the bound derived earlier by Borisov and Grishina using the same method. Applying a similar argument to magnetic white dwarf stars results in the more stringent bound $g_{ae}\\lesssim 9 \\cdot 10^{-13} (T/10^7 {\\rm K})^{5/4} (B/10^{10} {\\rm G})^{-2}$, where $T$ is the internal temperature of the white dwarf. ", + "introduction": "The axion $a$ is a pseudoscalar boson introduced by Peccei and Quinn (PQ) to solve the strong CP-problem in a natural way \\cite{pe77}. It results as a massless Goldstone boson from the spontaneous breaking of the PQ-symmetry by the VEV of a scalar field, $\\langle\\phi\\rangle = f_{\\rm PQ}$. As noted first by Weinberg and Wilczek \\cite{ww}, the chiral anomaly of QCD induces, at low temperatures $T\\lesssim\\Lambda_{\\rm QCD}\\ap 200\\:$MeV, a mass $m_a \\ap \\Lambda_{\\rm QCD}^2 / f_{\\rm PQ}$ for the axion. The numerical value of the axion mass is given by \\begin{equation} m_a = \\left( \\frac{ 0.60\\times 10^7 \\mbox{GeV}}{f_{\\rm PQ}} \\right) \\mbox{eV} \\:. \\end{equation} Besides the coupling to two photons through the chiral anomaly, the axion interacts with fermions by derivative couplings, which are all inversely proportional to the axion mass. The interaction of axions with electrons is described by \\begin{equation} \\label{Lae} {\\cal L}_{ae} = \\frac{g_{ae}}{2m_e}\\: \\bar\\psi\\gamma^\\mu\\gamma^5\\psi \\partial_\\mu a , \\end{equation} where the coupling constant \\begin{equation} g_{ae}=\\frac{m_e}{f_{\\rm PQ}} \\: c_e \\end{equation} depends on specific models through the effective PQ charge $c_e$. At tree level, the electron has $c_e= 1/3 \\cos\\beta$ in the DSFZ axion model, while it has $c_e =0$ in the KSVZ model \\cite{models}. In invisible axion models, the axion mass $m_a$ is in principle arbitrary, but in fact severely constrained by astrophysical and cosmological considerations \\cite{ko90,ra96}. The astrophysical bounds on $m_a$ arise because axion emission is an additional energy loss mechanism for stars, and therefore changes the stellar evolution. Since the axion couplings to matter are proportional to $m_a$, axion emission is suppressed for small values of $m_a$. Therefore astrophysics yields an upper bound for $m_a$. In ref. \\cite{ra95}, the bound $m_a \\lesssim 9 \\cdot 10^{-3}$eV was derived for a DFSZ axion by studying the evolution of red giant stars. On the other hand, cosmology provides a lower bound for $m_a$. Axions would have been produced in the early Universe as coherent waves or by the decay of axion strings. Therefore they could constitute cold dark matter. Requiring that axions do not overclose the Universe yields the lower bound $ 10^{-5}$eV $\\lesssim m_a$. It is the purpose of the present work to examine the emission of axions by electrons in the magnetic field $B$ of a strongly magnetized neutron star. The process considered, $e^- \\to e^- + a$, is very similar to cyclotron emission, $e^- \\to e^- + \\gamma$, and, therefore, may be called axion cyclotron emission. This process was first examined by Borisov and Grishina \\cite{bo94} using a semiclassical expression for the transition probability of the elementary process $e^- \\to e^- + a$, valid a priori only in the limit of low magnetic field strengths, $B\\ll B_{\\rm cr}= m_e^2 /e \\ap 4.41 \\times 10^{13}\\:$Gauss, and of ultrarelativistic electrons, $E\\gg m_e$. By comparing the axion cyclotron emission with neutrino cyclotron emission $e^- \\to e^- + \\nu + \\bar\\nu$, they concluded that the axion electron coupling constant is bounded from above by $g_{ae} \\lesssim 5 \\cdot 10^{-14}$. If this result, corresponding to $m_a \\lesssim 6 \\cdot 10^{-4}$eV for $c_e =1$, could be confirmed or improved, the window in the axion mass range would become very narrow, or even be closed. Consequently, it is imperative to recalculate the axion cyclotron emissivity and to extend the results obtained in ref. \\cite{bo94} to arbitrary magnetic field strengths and electron energies. We derive the probability of the elementary process $e^- \\to e^- + a$ taking the magnetic field exactly into account, and then numerically evaluate an expression for the axion emissivity which is exact in the limit of a completely degenerate Fermi-Dirac gas. As the main result of this work we obtain the axion cyclotron luminosity of a degenerate star with arbitrary core magnetic field strength $B$. Although we confirm the analytical approximations of ref. \\cite{bo94} in the limit $B\\ll B_{\\rm cr}$ and $E\\gg m_e$, the bound derived by us for the axion electron coupling constant applying the energy-loss argument to neutron stars is three order of magnitudes weaker. The main reason for this is that all conventional neutrino emission mechanisms such as the modified URCA process, which we take into account, were neglected in ref. \\cite{bo94}. % Finally, we consider the energy loss of white dwarf stars due to axion cyclotron emissivity. Requiring that for a white dwarf star the axion luminosity is smaller than the photon surface luminosity we obtain the bound $g_{ae} \\lesssim 9 \\cdot 10^{-13} (T/10^7{\\rm K})^{5/4} \\, (B/10^{10}{\\rm G})^{-2}$. ", + "conclusions": "We have derived the axion emissivity of a magnetized electron gas due to the process $e^- \\to e^- +a$ for arbitrary magnetic field strengths $B$. Comparing the exact emissivity with the semiclassical approximation obtained earlier, we have shown that this approximation is not only valid for $B/B_{\\rm cr} \\ll 1$, but also under the more general condition $E_F^2 \\gg 2eB$. Applying axion cyclotron emission as an additional cooling mechanism to neutron stars and requiring that the axion luminosity is smaller than the neutrino luminosity we could constrain the axion electron coupling constant to a value as small as $g_{ae} \\lesssim {\\cal O} (10^{-10})$. This bound is three orders of magnitude weaker than the bound derived in ref. \\cite{bo94} considering the same process. This discrepancy is not caused by differences in the calculation of the axion emissivity but by different assumptions about neutron star cooling. In the case of white dwarfs we could derive the more stringent limit eq. (\\ref{wd,limit}). We have noted that the lack of low-temperature magnetized white dwarfs could be interpreted as signature of an additional energy loss due to axion cyclotron emission." + }, + "9701/astro-ph9701182_arXiv.txt": { + "abstract": "Driftscan methods are highly efficient, stable techniques for conducting extragalactic surveys in the 21cm line of neutral hydrogen. Holding the telescope still while the beam scans the sky at the sidereal rate produces exceptionally stable spectral baselines, increased stability for RFI signals, and excellent diagnostic information about system performance. Data can be processed naturally and efficiently by grouping long sequences of spectra into an image format, thereby allowing thousands of individual spectra to be calibrated, inspected and manipulated as a single data structure with standard tools that already exist in astronomical software. The behavior of spectral standing waves (multi-path effects) can be appraised and excised in this environment, making observations possible while the Sun is up. The method is illustrated with survey data from Arecibo and Nan\\c{c}ay. ", + "introduction": "The designs of the Arecibo and Nan\\c{c}ay telescopes make them particularly well suited to taking data in a driftscan mode. Both have large collecting areas and are therefore sensitive survey instruments. They obtain the large collecting area by having much of their structure fixed to the ground. When they are used to track specific celestial coordinates, the on-axis gain changes and the far out sidelobes move in unpredictable ways, causing spill-over on the ground and RFI to be time variable. These instabilities increase the level of systematic uncertainties, thus increasing the difficulty of detecting weak signals. Driftscan observations avoid these problems since all components of the structures are fixed relative to the ground, thus achieving the full sensitivity of the large reflecting areas. Similar arguments apply to more conventional radio telescopes, when variation in the spillover causes fluctuations in the spectral baselines and when RFI entering the receiving system is modulated by gain variations in the far sidelobes as the antenna tracks celestial sources. This report is based on experience obtained in extragalactic HI surveys using the Arecibo and Nan\\c{c}ay Telescopes. Sorar (1994) and Briggs conducted the Arecibo HI Strip Survey in the driftscan mode in order to determine the HI-mass function for nearby extragalactic objects by surveying long strips at constant declination. The observations covered approximately 6000 independent sightlines to a depth of 7500~km~s$^{-1}$, but since the strips were retraced on many days in order to increase the integration time on each sightline, nearly a million individual spectra had to be calibrated, regrouped and averaged. The details of the observational technique were developed by Sorar (1994), and the results and followup to the survey are summarized by Zwaan et al (1996, these proceedings). \\begin{figure} \\centerline{ } \\caption{Nan\\c{c}ay Raw Data Image. Passband calibrated spectra have been loaded into the image in time sequence increasing from right to left. There are two slightly overlapping, 6.4 MHz wide spectral bands, as marked at the right border. A trace of the continuum as a function of right ascension is drawn under the image. } \\end{figure} The method is also in use for surveys at Nan\\c{c}ay. One project, being conducted by Kraan-Korteweg, van Driel, Binggeli, and Briggs, will test the completeness to a depth of 2300 km~s$^{-1}$ of deep optical catalogs of dwarfs and low surface brightness galaxies in the CnV~I cloud (Binggeli et al 1990). A sample of the raw data from the HI survey are shown in Figure 1. The spectral passband calibration for this data has been performed by computing the average spectrum of the entire 3.5 hour dataset, and each spectrum was divided by the average spectrum as it was loaded in time sequence into the columns of the image. Passage of a telescope beam over a background continuum source is registered in the image as a dark band. The residual Galactic HI emission causes the splotchy horizontal band across the image around the HI rest frequency. Once the data is loaded in image format, images become the units in which the data is stored, manipulated and displayed. For example, continuum subtraction, averaging of data from different days and smoothing can be accomplished using procedures in familiar astronomical image processing packages. Figure 2 shows the processed spectra for a total of $\\sim$2500 sightlines in three adjacent declination strips. The figure results from approximately 82,000 individual spectra. There are 11 detected extragalactic signals resulting from 9 separate galaxies, plus a number of interesting features associated with Galactic HI. A second project now in progress at Nan\\c{c}ay will observe ${\\sim}6$\\% of the sky to a depth of 4500~km~s$^{-1}$ with noise level ${\\sim}23$ mJy (5$\\sigma$) for velocity resolution 20 km~s$^{-1}$. The survey is well matched to detecting nearby examples of gas-rich systems such as HI~1225+01 (Giovanelli \\& Haynes 1989, Chengalur, Giovanelli \\& Haynes 1995) and the circum-galactic ring in Leo (Schneider 1989), as well as detecting a sample of several hundred normal galaxies. \\begin{figure} \\centerline{ } \\caption{Processed Nan\\c{c}ay images for three declinations. Spectra are loaded in horizontal lines in this image, with right ascension increasing upwards. Overlap of the spectral sub-bands has been removed, and labels on the horizontal axis indicate velocity in km~s$^{-1}$. Eleven detected signals are marked.} \\end{figure} ", + "conclusions": "" + }, + "9701/astro-ph9701077_arXiv.txt": { + "abstract": "We have systematically investigated the {\\it ASCA} spectra of 12 early type galaxies. This paper presents the global spectral properties of these systems based on a larger sample than in any previous {\\it ASCA} study. The X-ray spectra were uniformly fitted by a two-component model consisting of hard X-rays from thermal emission with a temperature of about 10 keV or from a power-law with index 1.8, plus soft X-rays from a thin thermal plasma with temperature ranging from 0.3 to 1 keV. The X-ray luminosities of the hard component are found to be proportional to the blue band luminosities, while those of the soft component show large scatter with no clear correlation. The metal abundances determined from the soft component are systematically lower than solar, with a mean value of about 0.3 solar. We examine the relationships between the temperature and volume emission measure, and between the gas temperature and the stellar velocity dispersion. The volume emission measures for early type galaxies plotted as a function of the gas temperature are well below the extrapolated line found in clusters of galaxies, indicating that early type galaxies are relatively gas poor compared with galaxy clusters. The ratio of the stellar kinetic energy per unit mass to the thermal energy of the hot gas per unit mass ($\\beta_{spec}$) is less than unity. We found no systematic relationship between X-ray properties and environment, suggesting that the interaction between interstellar matter and the intracluster medium is not strong. ", + "introduction": "The {\\it Einstein} satellite discovered X-ray emitting gas halos around luminous early type galaxies (eg. Forman, Jones, \\& Tucker 1985 \\markcite{FJT85}), forcing the rejection of previous assumptions that elliptical galaxies are gas poor systems. The X-ray surface brightness distribution was found to closely follow the optical image (Trinchieri, Fabbiano, \\& Canizares \\markcite{TFC86} 1986); however the X-ray luminosity, and hence the mass of the hot gas, showed large scatter from galaxy to galaxy (eg. Canizares, Fabbiano, \\& Trinchieri \\markcite{CFT87} 1987). The estimated cooling time of the hot gas is less than the Hubble time (Trinchieri et al. \\markcite{TFC86} 1986), however cooling may be balanced by heating mechanisms such as supernova explosions. Arnaud et al. \\markcite{Arn92} (1992) found that the iron mass in the intracluster medium (ICM) is directly proportional to the optical luminosity from early type galaxies in clusters of galaxies. This leads one to suspect that metals in the ICM originate in early type galaxies and that there might be a link, in the form of early star formation and galactic winds, between the ICM and metal abundances in early type galaxies. Due to the poor energy resolution of previous satellites, however, critical parameters for investigating the nature of the hot gas in early type galaxies, such as the gas temperature, mass, and metal abundance have not been well-constrained. Accordingly, several important issues including the origin, heating, and fate of the hot gas, and the effect of early type galaxies on the ICM, are not well understood. {\\it ASCA} (Tanaka et al. \\markcite{TIH94} 1994) is the first satellite with sufficient energy resolution and large effective area over the broad 0.4 to 10 keV bandpass to determine these physical parameters. With this satellite, Awaki et al. \\markcite{Awaki94} (1994) analyzed 3 early type galaxies and found that the temperature and metal abundance of the hot gas are about 0.8 keV and 0.4 solar, respectively. The inferred metal abundance is lower than expected from theoretical models, and lower than those determined at other wavelengths. Accordingly, the {\\it ASCA} observations raise a serious challenge to standard models for the evolution of galaxies and clusters of galaxies. This mystery deepened with the discovery of even lower metallicities (less than 0.2 solar) in two other elliptical galaxies (Loewenstein et al. \\markcite{LMT94} 1994), and in the outer regions of NGC4636 (Mushotzky et al. \\markcite{MLA94} 1994). Matsushita et al. \\markcite{Kyoko94} (1994) discovered, analyzing 5 early type galaxies including the sample of 3 in Awaki et al. \\markcite{Awaki94}(1994), a hard X-ray component at energies above a few keV in addition to the hot gas component dominating below about 2 keV. In order to perform plasma diagnostics of the hot gas component, it is essentially important to estimate and remove the hard component from the X-ray spectra of early type galaxies. If the hard component is due to the integrated emission of stellar sources, as suggested by Matsushita et al. \\markcite{Kyoko94} (1994), its luminosity should be proportional to optical luminosity. Since the samples in previous {\\it ASCA} studies are too limited to fully investigate all the issues mentioned above, we have systematically studied twelve early type galaxies, including the samples of Awaki et al. \\markcite{Awaki94}(1994) and Matsushita et al. \\markcite{Kyoko94} (1994) as well as newly observed galaxies. These comprise the largest sample of {\\it ASCA} spectra of early type galaxies yet investigated. ", + "conclusions": "We have presented results from the uniform analysis of the {\\it ASCA} spectra of 12 early type galaxies. This is the largest such sample to date and, as summarized below, we have uncovered new facts as well as confirmed previous inferences based on much smaller samples. \\begin{enumerate} \\item We found that the X-ray spectra from the 12 galaxies in our sample generally consist of at least two components: a soft component from X-ray emission of hot gas, and a hard component that can be characterized by either a thermal bremsstrahlung model with $kT>5$ keV or a power-law model of photon index 1.8. The consistency of the strength and spectral shape of the hard component with those in bulge-dominated spirals corroborates previous suggestions that it primarily originates as the integrated emission from X-ray binaries. The hard component spectrum requires no significant Fe K line emission, and the upper limit on the equivalent width of Fe K line emission is 460 eV. \\item We confirmed that the metal abundance of the hot gas is generally less than 0.5 solar. There is no clear correlation between the abundance and temperature of the hot gas, however the lowest abundances are found in the coolest galaxies. \\item There is a large dispersion in the ratio of gas mass ({\\it i.e.} the X-ray luminosity of the soft component) to optical luminosity, implying a wide variety of gas-dynamical histories. \\item Theories of large-scale structure imply that there is a continuum of dark matter halos spanning galaxy and cluster scales. However, we have found that early type galaxies fall well below the constant gas mass fraction extension of the hottest clusters in the $kT-VEM$ plane (VEM $\\propto kT^{1.5}$), and that the ratio of thermal to kinetic energy per unit mass is, on average, twice that found for clusters of galaxies. This suggests that larger systems are generally more efficient at retaining their hot gas and less efficient at converting gas into stars. The interstellar media of elliptical galaxies are not simply scaled down version of intracluster media, but have been more profoundly effected by feedback from star formation and other gas-dynamical processes. \\item The low Fe abundances, as well as the absence of a negative correlation between the ``excess temperature'', defined as the difference between the hot gas temperature and the temperature corresponding to the stellar velocity dispersion, and the ratio of the square root of the volume emission measure of the hot gas to the blue band luminosity (an indicator of the mass ratio of hot gas to stars) argue against type Ia supernovae being an important heat source for the ISM. \\item We found no systematic differences in the physical parameters of the hot gas between galaxies in dense and tenuous intracluster environments, although our ability to test this is limited. This indicates that interaction between the intracluster medium and interstellar matter may not be strong. \\end{enumerate} In short, {\\it ASCA} analysis of early type galaxies has yielded both expected (a hard component with the same normalization as in spiral bulges) and unexpected (low interstellar abundances) results. The interrelationships and dispersion among the properties of the hot gas component, and their comparison with those in intracluster media demonstrates the relative importance of astration, mass loss, and gas-dynamics in determining the evolution of early type galaxies and their ISM. Any complete theory of galaxy formation must explain the observed trends and diversity of bound distributions of hot gas on galaxy and cluster scales. \\vspace*{1em}" + }, + "9701/astro-ph9701196_arXiv.txt": { + "abstract": "We present a statistical comparison of three different estimates of cluster mass, namely, the dynamical masses obtained from the velocity dispersion of optical galaxies, the X-ray masses measured from the temperature of X-ray emitting gas under the assumption of isothermal hydrostatic equilibrium, and the gravitational lensing masses derived from the strong/weak distortions of background galaxy images. Using a sample of 29 lensing clusters available in literature, we have shown that the dynamical masses are in agreement with the gravitational lensing masses, while the X-ray method has systematically underestimated cluster masses by a factor 2-3 as compared with the others. These results imply that galaxies indeed trace the gravitational potential of their clusters, and there is no bias between the velocities of the dark matter particles and the galaxies in clusters. The X-ray cluster mass discrepancy is probably from the simplification in the models for the X-ray gas distribution and dynamical evolution. ", + "introduction": "Clusters of galaxies are the largest coherent and gravitationally bounded objects in the universe. The precise determination of their gravitational masses is crucial for our understanding of formation and evolution of cosmic structures, for mapping of matter distribution on large-scales and also for measurement of the present mean mass density of the universe ($\\Omega_0$). Historically, cluster masses are derived from the dynamical analysis of the observed velocity dispersion of cluster galaxies based on the virial theorem, which results in the so-called virial cluster masses $M_v$. With the development of the X-ray astronomical techniques clusters can be selected from the X-ray emission of the hot diffuse intracluster gas, giving rise to the X-ray cluster masses $M_x$ when combined with the assumption of hydrostatic equilibrium. Over the past decade the detection of the gravitationally distorted images of faint distant galaxies behind some clusters of galaxies provides another independent mass estimate: the gravitational cluster masses $M_{lens}$. In particular, $M_{lens}$ are obtained regardless of the cluster matter state and components. These three methods should be incorporate, and comparisons of their results would yield a very useful clue regarding the dynamical evolution of clusters and a test for the accuracy of cluster mass determinations. The early studies based on a few selected clusters in which both lensing and optical and/or X-ray data are available claimed a cluster mass discrepancy by a factor of typically $2\\sim3$ among the three methods (Wu, 1994; Fahlman et al. 1994; Miralda-Escud\\'e \\& Loeb 1995), while a consistency between dynamical masses and gravitational lensing derived ones has been also reported in some cases, e.g. PKS0745-191 (Allen, Fabian \\& Kneib 1996). We have recently carried out a statistical comparison of the overall cluster radial matter distributions determined from X-ray observation and gravitational lensing and concluded that the X-ray analysis may have systematically underestimated cluster masses at least in the central regions (Wu \\& Fang 1996; hereafter Paper I). From the theoretical point of view, the above three methods have to give the same cluster masses if clusters are dynamically relaxed. This probably accounts for the result in clusters like PKS0745-191 which appear to be regular in optical/X-ray morphologies. Though one may attribute the reported mass discrepancy to the nonthermal pressure (Loeb \\& Mao 1994; Ensslin et al. 1997) or the projection effect (Miralda-Escud\\'e \\& Babul 1995; Cen 1996), it is most likely that the problem is relevant to the cluster matter distribution and dynamical evolution. Indeed, both optical and X-ray observations have revealed the presence of substructures in most of clusters, implying that clusters may be still in the era of formation. In particular, the recent spatially-resolved measurements of gas temperature in some clusters illustrate the complex two-dimensional patterns including the asymmetric variations and the significant decline with radius (e.g. Henriksen \\& White 1996; Henriksen \\& Markevitch 1996; Markevitch 1996), which are the strong indicators of the effect of substructure merging. For those clusters, we are unable to apply the equation of hydrostatic equilibrium to the X-ray emitting gas. Therefore, it is naturally expected that the X-ray cluster mass obtained under the assumption of a hydrostatic equilibrium is in principle not representative of the true cluster mass, and thereby should be different from the gravitational lensing derived mass and/or even the virial mass. This scenario is supported by the recent numerical study of X-ray and lensing properties of clusters of galaxies from the standard CDM simulations by Bartelmann \\& Steinmetz (1996). They found that the cluster masses can be biased low if the traditional $\\beta$ model is adopted for the intracluster gas, which is especially true when clusters exhibit pronounced substructures, while the strong lensing preferentially selects the clusters that are dynamically more active than the average. Alternatively, it should be noticed that $M_v$ derived from the virial theorem and $M_x$ derived from the hydrostatic equilibrium have very different physical implications. $M_v$ is related to the galaxy velocity dispersion and galaxy number density. Comparison of $M_v$ and $M_{lens}$ would set constraints on the bias parameter between the velocities of the galaxies and of the dark matter and test whether optical galaxies trace the gravitational potential of the cluster. On the other hand, $M_x$ depends on the temperature variation and the density profile of the hot diffuse gas, which may suffer from the influence of the possible existence of the turbulence and magnetic field. As a result, comparison of $M_x$ and $M_v$ may allow one to determine how significant the nonthermal pressure contributes to the computation of $M_x$ because galaxies are unaffected by the nonthermal pressure in clusters. Furthermore, the comparison of the three cluster estimates may help to solve the puzzle of the baryonic matter excess in clusters if clusters provide a fair sample of the universal baryon fraction (see Paper I for summary). Recall that the baryon fractions in clusters are computed using the dynamical masses $M_x$ and/or $M_v$, which may have large uncertainties if most of clusters are still in the process of violent merging. As a whole, galaxies and gas particles have probably experienced very different evolutions in the formation of clusters and then exhibit different dynamical states and density distributions in clusters today. Numerous lensing and optical/X-ray observations of clusters have now made it timely and possible to carry out these comparisons. Unlike Paper I that chose X-ray and lensing data separately from literature, we now work with the lensing clusters only, in which the strongly and/or weakly distorted images of background galaxies have been observed. While this paper was in the refereeing stage, we received a preprint by Smail et al. (1997) who made a similar but sophisticated investigation for 11 distant clusters observed with {\\it HST}. Throughout this paper we adopt a matter-dominated flat cosmological model of $\\Omega_0=1$ and a Hubble constant of $H_0=50$ km s$^{-1}$ Mpc. ", + "conclusions": "Comparisons of cluster mass estimates from X-ray gas hydrostatic equilibrium, dynamical analysis and gravitational lensing are shown in Figure 1(a) and (b) using the data of Table 1. An immediate conclusion is that $m_{lens}$ agrees essentially with $m_v$ while a systematic excess of $m_{lens}$ with respect to $m_x$ is detected. This can be clearly demonstrated by the following best-fit relation to the data: \\begin{equation} m_{lens}=(1.42\\pm0.99)m_v =(2.23\\pm1.15)\\beta_{fit}^{-1}m_x, \\end{equation} in which the uncertainties are the 1$\\sigma$ errorbars. Meanwhile, Fig.1(a) and (b) also illustrate the influence of cluster morphologies on the relations between $m_{lens}$ and $m_v$ and between $m_{lens}$ and $m_x$, respectively. Apparently, it is very unlikely that cluster morphologies can lead to a remarkable difference in the results. \\placefigure{fig1} In Fig.2(a) and (b) we display the variations of $m_{lens}$, $m_v$ and $m_x$ with the cluster radius. Basically, Fig.2(b) provides a result similar to the one of Paper I, i.e., there is a systematic discrepancy between $m_x$ and $m_{lens}$ inside the cluster core radius, and the projected gravitational cluster mass obtained with lensing mothed follows a simple power-law of $\\sim r^{1.3}$. Again, in contrast to $m_x$, $m_v$ agrees statistically with $m_{lens}$ over all the scales. \\placefigure{fig2} Additionally, we have computed the $\\beta$ parameter characterizing the specific energies of the galaxies and the gas in clusters, $\\beta_{spec}\\equiv \\sigma^2/(kT/\\mu m_p)$ where $\\mu m_p=0.59$ is the mean particle mass. Our best-fit value with the 18 data points (both $T$ and $\\sigma$ are available) in Table 1 reads $\\beta_{spec}=1.29\\pm0.71$, while $\\beta_{spec}$ reduces to $1.17\\pm0.50$ if AC114 is excluded. The best-fit relation between the galaxy velocity dispersion and the gas temperature is \\begin{equation} (\\sigma/{\\rm km\\; s}^{-1})=10^{2.64\\pm0.11}(T/{\\rm keV})^{0.51\\pm0.13}. \\end{equation} It appears that our best-fit average $\\beta_{spec}$ and $\\sigma$--$T$ relation are consistent with the previous work [see Girardi et al. (1996) for summary]. Based on such a good fitness of eq.(5) alone, one might conclude that the galaxies and the gas are in hydrostatic equilibrium with the same cluster potential, as was claimed by Lubin \\& Bahcall (1993). However, our result of eq.(4) raises a new question of how one could reconcile the discrepancy between $m_{lens}$ (or $m_v$) and $m_x$ with the good correlation between $\\sigma$ and $T$. Alternatively, it turns out from Fig.1(a) that a zero core radius $r_{dc}=0$ for the cluster mass profile provides a good fit to the lensing data. The best-fit core radius $r_{dc}$ by requiring $m_{v}=m_{lens}$ is $r_{dc}=-0.09\\pm0.24$ Mpc, indicative of rather a compact dark matter distribution in clusters. This result is compatible with the early studies of giant arcs and statistical lensing of arcs/arclets which report a small core radius ($<0.1$ Mpc) for the dark matter profile of the arc clusters (e.g. Hammer 1991; Wu \\& Hammer 1993; Grossman \\& Saha 1994). While there is a significant evolution of X-ray luminosity clusters with redshift (Edge et al. 1990; Gioia et al. 1990; Henry et al. 1992), the deficit of the X-ray cluster mass may be relevant to the cluster evolution. Recall that the local $L_x$--$T$ relation established at low redshift $z<0.1$ was employed to estimate the cluster temperature, whereas most of the clusters in Table 1 are actually located at intermediate redshift $z\\sim0.2$--$0.5$. Therefore, it would be useful to examine the dependence of the ratio of $m_{lens}$ to $m_x$ on the cluster redshift. Our best-fit relation is \\begin{equation} \\frac{m_{lens}}{m_x}=10^{0.09\\pm0.10}(1+z)^{1.7\\pm0.9} \\end{equation} Namely, the cluster mass discrepancy between the X-ray analysis and the gravitational lensing method is indeed related to the evolutionary history of clusters. Since $m_x$ is proportional to $T$ according to eq.(1), the cluster temperature has a similar variation with redshift. This yields a temperature ratio of $1.5_{-0.3}^{+0.4}$ for cluster at redshift $z=0.33$ relative to the one at $z=0.035$, in consistent with the result ($1.4_{-0.3}^{+0.4}$) of Henry et al. (1994). Though the cluster temperature evolution since intermediate redshift is moderate, it may account for the mass discrepancy we report in the present paper. We now discuss briefly the significance of the consistency/discrepancy between the cluster masses derived from gravitational lensing, dynamical analysis and X-ray observations. Both gravitational lensing and ``virial'' methods yield nearly the same cluster masses, which have several implications: First, galaxy velocity dispersion indeed provides a good estimate of cluster mass. Second, the dark matter particles and the galaxies have approximately the same velocity dispersion, i.e., there is no velocity bias in clusters of galaxies. Third, mass follows the light. These arguments are comparable with the recent dynamical analysis of the CNOC cluster sample (Carlberg et al. 1996), which has found strong evidence that galaxies are effectively in equilibrium with their host cluster. However, our finding disagrees with the numerical result that the velocity biasing parameter is $\\sim0.7$--$0.8$ (Carlberg \\& Dubinski 1991; Couchman \\& Carlberg 1992). On the other hands, cluster mass estimate based on the X-ray temperature assuming an isothermal hydrostatic equilibrium has systematically underestimated cluster masses by a factor of $\\sim2$--$3$ as compared with gravitational lensing and ``virial'' method, which demonstrates that the gas particles may not be a good tracer of the gravitational potential of the cluster. The recent high-spectral resolution observations do reveal the occurrence of complex temperature maps, indicating the on-going merger activities in clusters and the cluster evolution with cosmic epoch. Overall, the simplification of modeling the temperature (isothermal) and gas density profile (spherical) is responsible for the deficit of the X-ray cluster mass detected in this paper. One may also attribute the mass disprepancy to the nonthermal pressure in clusters (Loeb \\& Mao 1994; Ensslin et al. 1997), which affects the gas particles while produces no effect on galaxy distribution and velocity dispersion. Finally, it is pointed out that our results still contain large scatters due to the scarcity of the lensing data, and a large cluster sample will be needed to confirm our finding." + }, + "9701/astro-ph9701019_arXiv.txt": { + "abstract": "We present a comprehensive hybrid library of synthetic stellar spectra based on three original grids of model atmosphere spectra by Kurucz (1995), Fluks {\\em et al.\\/} (1994), and Bessell {\\em et al.\\/} (1989, 1991), respectively. The combined library has been intended for multiple-purpose synthetic photometry applications and was constructed according to the precepts adopted by Buser \\& Kurucz (1992): (i) to cover the largest possible ranges in stellar parameters (${\\rm T_{eff}}$, log g, and [M/H]); (ii) to provide flux spectra with useful resolution on the uniform grid of wavelengths adopted by Kurucz (1995); and (iii) to provide synthetic broad--band colors which are highly realistic for the largest possible parameter and wavelength ranges. Because the most astrophysically relevant step consists in establishing a {\\em realistic\\/} library, the corresponding color calibration is described in some detail. Basically, for each value of the effective temperature and for each wavelength, we calculate the {\\em correction function\\/} that must be applied to a (theoretical) solar--abundance model flux spectrum in order for this to yield synthetic UBVRIJHKL colors matching the (empirical) color--temperature calibrations derived from observations. In this way, the most important systematic differences existing between the original model spectra and the observations can indeed be eliminated. On the other hand, synthetic UBV and Washington ultraviolet excesses $\\delta_{(U-B)}$ and $\\delta_{(C-M)}$ and $\\delta_{(C-T_{1})}$ obtained from the original giant and dwarf model spectra are in excellent accord with empirical metal--abundance calibrations (Lejeune \\& Buser 1996). Therefore, the calibration algorithm is designed in such a way as to {\\em preserve the original differential grid properties implied by metallicity and/or luminosity changes\\/} in the new library, if the above correction function for a solar--abundance model of a given effective temperature is also applied to models of the same temperature but different chemical compositions [M/H] and/or surface gravities log g. While the new library constitutes a first--order approximation to the program set out above, it will be allowed to develop toward the more ambitious goal of matching the full requirements imposed on a {\\em standard library\\/}. Major input for refinement and completion is expected from the extensive tests now being made in population and evolutionary synthesis studies of the integrated light of globular clusters (Lejeune 1997) and galaxies (Bruzual {\\em et al.\\/} 1997). ", + "introduction": "The success of population and evolutionary synthesis calculations of the integrated light of clusters and galaxies critically depends on the availability of a suitable library of stellar spectral energy distributions (SEDs), which we shall henceforth call {\\em stellar library\\/}. Because of the complex nature of the subject, there are very many ways in which such calculations can contribute to the solution of any particular question relevant to the stellar populations and their evolution in clusters and galaxies. In previous studies, the particularities of these questions have largely determined the properties that the corresponding stellar library must have in order to be considered {\\em suitable\\/} for the purpose. There is now a considerable arsenal of observed stellar libraries (for a recent compilation, see e.g. Leitherer {\\em et al.\\/} 1996), each of which has its particular resolution, coverage, and range of wavelengths as well as its particular coverage and range of stellar parameters -- but which, even if taken in the aggregate, fall short by far of providing the {\\em uniform, homogeneous, and complete\\/} stellar library which is required now for a more systematic and penetrating exploitation of {\\em photometric\\/} population and evolutionary synthesis. Ultimately for this purpose, what is needed is a uniform, homogeneous, and complete {\\em theoretical\\/} stellar library, providing SEDs in terms of physical parameters consistent with empirical calibrations at all accessible wavelengths. Thus, the above goal can be approached by merging existing grids of theoretical model--atmosphere spectra into the desired uniform and complete stellar library, and making it both homogeneous and realistic by {\\em empirical calibration\\/}. A variant of this approach was first tried by Buser \\& Kurucz (1992), who constructed a more complete theoretical stellar library for O through K stars by merging the O--G--star grids of Kurucz (1979a,b) with the grids of Gustafsson {\\em et al.\\/} (1975), Bell {\\em et al.\\/} (1976), and Eriksson {\\em et al.\\/} (1979) for F--K stars. In their paper, Buser \\& Kurucz solved for uniformity and homogeneity by recomputing new late--type spectra for Kurucz's (1979) standard grid of wavelengths and using the Kurucz \\& Peytremann (1975) atomic opacity source tables. The resulting hybrid library\\footnote{The Buser--Kurucz library is a {\\em hybrid\\/} library in the sense that it is based on {\\em two distinct grids of model atmospheres\\/} which were calculated using different codes (ATLAS AND MARCS, respectively); it is {\\em quasi--homogeneous\\/}, however, because for both grids of model atmospheres the {\\em spectrum calculations\\/} were obtained using the {\\em same opacity source tables\\/}.} has, indeed, significantly expanded the ranges of stellar parameters and wavelengths for which {\\em synthetic photometry\\/} can be obtained with useful systematic accuracy and consistent with essential empirical effective temperature and metallicity calibrations (Buser \\& Fenkart 1990, Buser \\& Kurucz 1992, Lejeune \\& Buser 1996). In the meantime, Kurucz (1992, 1995) has provided a highly comprehensive library of theoretical stellar SEDs which is homogeneously based on the single extended grid of model atmospheres for O to late--K stars calculated from the latest version of his ATLAS code and using his recent multi--million atomic and molecular line lists. The new Kurucz grid -- as we shall call it henceforth -- indeed goes a long way toward the {\\em complete library\\/} matching the basic requirements imposed by synthetic photometry studies in population and evolutionary synthesis. As summarized in Table 1 below, SEDs are provided for uniform grids of wavelengths and stellar parameters with almost complete coverage of their observed ranges ! These data have already been widely used by the astronomical community, and they will doubtlessly continue to prove an indispensable database for population and evolutionary synthesis work for years to come. In the present work, we shall endeavor to provide yet another indispensable step toward a {\\em more complete\\/} stellar library by extending the new Kurucz grid to {\\em cooler temperatures\\/}. This extension is particularly important for the synthesis of old stellar populations, where {\\em cool giants and supergiants\\/} may contribute a considerable fraction of the total integrated light. Because model atmospheres and flux spectra for such stars -- the M stars -- have been specifically calculated by Bessell {\\em et al.\\/} (1989, 1991) and by Fluks {\\em et al.\\/} (1994), our task will mainly be to combine these with the new Kurucz grid by transformation to the same uniform set of wavelengths, and to submit the resulting library to extensive tests for its realism. In fact, as shall be shown below, the process will provide a complete grid of SEDs which is homogeneously and consistently {\\em calibrated against observed colors\\/} at most accessible wavelengths. In Sect. 2, we shall briefly describe the different libraries used in this paper and the main problems that they pose to their unification. Because the spectra exhibit systematic differences both between their parent libraries and relative to observations, we set up, in Sect. 3, the basic empirical color--temperature relations to be used for uniformly calibrating the library spectra in a wide range of broad--band colors. This calibration process is driven by a computer algorithm developed and described in Sect. 4. The actual color--temperature relations obtained from the corrected library spectra are discussed in Sect. 5, and the final organization of the library grid is presented in Sect. 6. In the concluding Sect. 7, we summarize the present state of this work toward the intended {\\em standard library\\/}, and we briefly mention those necessary steps which are currently in process to this end. ", + "conclusions": "Although astronomers have for a long time agreed that a uniform, complete, and realistic stellar library is urgently needed, it must be emphasized that this goal has remained too ambitious to be achieved in a single concerted effort through the present epoch. We have thus attempted to proceed in well--defined steps, with priorities set according to the availability of basic data and following the most obvious scientific questions that would likely become more tractable, or even answerable. Therefore, we briefly review the present achievement to clarify its status in the ongoing process toward a future {\\em standard library\\/} of theoretical stellar spectra for photometric evolutionary synthesis. \\begin{enumerate} \\item {\\bf Completeness.} The unification of the massive K--library spectra with those for high-- luminosity M--star models (the B+F--library) is most important, because even as a small minority of the number population of a given stellar system, the late--type giants and supergiants may provide a large fraction of this system's integrated light at visible and infrared wavelengths. This fact was recognized early on (e.g., Baum 1959), and eventually also co--motivated the effort leading to the existence of the B--library used in this work (Bessell {\\em et al.\\/} 1988). While the K--, B--, and F--libraries have been used to remedy incompleteness (in either wavelength or parameter coverage, or both) of available {\\em observed\\/} stellar libraries before (e.g., Worthey 1992, 1994; Fioc \\& Rocca--Volmerange 1996), our first goal here has been to join them as a purely {\\em theoretical\\/} library, providing the main advantages of physical homogeneity and definition in terms of fundamental stellar parameters -- which allows direct use with stellar evolutionary calculations. But even so, the present library remains incomplete in several respects. First, stellar evolution calculations (e.g., Green {\\em et al.\\/} 1987) predict that high--luminosity stars with temperatures near or below $3500K$ may also exist at low metallicities, $[M/H] \\sim -1.0$, and their flux contributions at visible--near ultraviolet wavelengths (where metallicity produces significant effects) may not quite be negligible in the integrated light of old stellar populations. Therefore, in order to provide the fuller coverage required for an adequate study of this metallicity--sensitive domain, new calculations of B--library spectra extending the original data to both shorter wavelengths $(\\lambda\\lambda \\geq 320nm)$ and lower abundances $([M/H] \\geq -2)$ (Buser {\\em et al.\\/} 1997) will replace the current hybrid B+F spectra and make the next library version more homogeneous. Secondly, even though the low--temperature, low--luminosity \\/ M--dwarf stars do not contribute significantly to the integrated bolometric flux, they are not negligible in the determination of mass--to--light ratios in stellar populations. Thus, a suitable grid of (theoretical) M--dwarf spectra calculated by Allard \\& Hauschildt (1995) is being subjected to a similar calibration process (Paper II) and will be implemented in the present library as an important step toward the intended {\\em standard stellar library\\/}. Finally, the libraries of synthetic spectra for hot O-- and WR--stars which were recently calculated by Schmutz {\\em et al.\\/} (1992) will allow us to extend the calibration algorithm to ultraviolet IUE colors, where such stars radiate most of their light. \\\\ \\item {\\bf Realism.} In view of its major intended application -- photometric evolutionary synthesis --, the {\\em minimum requirement\\/} that we insist the theoretical library {\\em must\\/} satisfy, is to provide stellar flux spectra having (synthetic) colors which are systematically {\\em consistent with calibrations derived from observations\\/}. How else could we hope to learn the physics of distant stellar populations from their integrated colors, unless the basic building blocks -- i.e., the library spectra used in the synthesis calculations -- can be taken as adequate representations of the better--known fundamental stellar properties, such as their color--temperature relations? Because the original library spectra do not meet the above minimum requirement (Sect. 3), we have developed an algorithm for calibrating existing theoretical spectra against empirical color--temperature relations (Sect. 4). Because comprehensive empirical data are unavailable for large segments of the parameter space covered by the theoretical library, direct calibration can be effected only for the major sequences of solar--abundance models (Sect. 5). However, we have also shown that the present algorithm provides the desired broad--band (or pseudo--continuum) color calibration {\\em without\\/} destroying the original relative monochromatic fluxes between arbitrary model grid spectra and solar--abundance calibration sequence spectra of the same effective temperature. This conservation of original grid properties also propagates with useful systematic accuracy even through most {\\em differential broad--band colors\\/} of the corrected library spectra. Thus, to the extent that differential broad--band colors of original library spectra were previously shown to be consistent with spectroscopic or other empirical calibrations of the UBV--, RGU--, and Washington ultraviolet-excess--metallicity relations (Buser \\& Fenkart 1990, Buser \\& Kurucz 1992, Lejeune \\& Buser 1996), the {\\em corrected} library spectra are still consistent with the same calibrations. \\\\ \\item{\\bf Library development.} At this point, we feel that some of the more important intrinsic properties required of the future {\\em standard library\\/} have already been established. Of course, many more consistency tests and calibrations will now be needed that can, however, only be performed for {\\em local\\/} volumes of the full parameter space covered by the new library. For example, we shall use libraries of observed flux spectra for individual field and cluster stars to better assess -- and/or improve -- the performance of the present library version in the non--solar--abundance and non--visual wavelength regimes. Eventually, we also expect significant guidance toward a more {\\em systematically\\/} realistic version of the library from actual evolutionary synthesis calculations of the integrated spectra and colors of globular clusters (Lejeune 1997). Last, but not least, we would like to emphasize that, while we here present the results taylored according to the general needs in the field of photometric evolutionary synthesis, the library construction algorithm has been designed such as to {\\em allow flexible adaptation to alternative calibration data as well\\/}. As we shall ourselves peruse this flexibility to accommodate both feed--back and new data, the reader, too, is invited to define his or her own preferred calibration constraints and have the algorithm adapted to perform accordingly. \\end{enumerate} {\\em Acknowledgements}. We are grateful to Michael Scholz and Gustavo Bruzual for providing vital input and critical discussions. Christophe Pichon is also aknowledged for his precious help with the final implementation of some of the figures. We wish to thank warmly the referee for his helpful comments and suggestions. This work was supported by the Swiss National Science Foundation." + }, + "9701/astro-ph9701006_arXiv.txt": { + "abstract": "Turbulent viscosity is believed to circularize and synchronize binary orbits and to damp stellar oscillations. It is also believed that when the tidal period is shorter than the turnover time of the largest eddies, turbulent viscosity is partially suppressed. The degree of suppression, however, is disputed. We re-examine both of these beliefs via (i) direct perturbative calculations, linearizing the fluid equations on a turbulent background; and (ii) numerical integration of a chaotic dynamical system subject to periodic forcing. We find that dissipation of rapid tides is severely suppressed. Furthermore, circularization of late-type binaries does not occur by turbulent convection if it occurs on the main sequence. \\medskip\\noindent {\\it POPe-703. Submitted to the Astrophysical Journal 12/19/96} ", + "introduction": "Eddy viscosity in convection zones is presumed to dissipate the shear associated with variable tides, and hence gradually to circularize and synchronize close binaries [\\cite{Z:66}, \\cite{Z:77}, \\cite{Z:92} and references therein]. This theory of tidal circularization has been tested against binaries containing giant stars, and the agreement is satisfactory (\\cite{VP:96}). This paper will focus on a theoretical controversy concerning eddy viscosity whose outcome may have consequences not only for tidal circularization, but also for stellar pulsations and oscillations. Assuming a Kolmogorov cascade of eddies, we have on scales $\\lambda$ smaller than the mixing length ($l$ that \\begin{eqnarray} v_\\lambda &\\approx& \\left(\\frac{\\lambda}{l}\\right)^{1/3} v_l, \\nonumber\\\\ \\tau_\\lambda &\\approx& \\left(\\frac{\\lambda}{l}\\right)^{-2/3} \\tau_l, \\nonumber\\\\ \\nu_\\lambda &\\approx& \\left(\\frac{\\lambda}{l}\\right)^{2/3} \\nu_l, \\label{eq:Kolmogorov} \\end{eqnarray} where $v_l$, $\\tau_l$, and $\\nu_l$ are the convective velocity, turnover time, and effective viscosity on the scale of the mixing length, which is comparabale to the pressure scale height, $H$; while $v_\\lambda$, $\\tau_\\lambda$, and $\\nu_\\lambda$ are the corresponding quantities on smaller scales $\\lambda < l$. If the eddies on scales are space-filling, then the total turbulent viscosity $\\nu_T\\approx\\nu_l$. But when the period ($P$) of the tide or oscillation is less than $\\tau_l$, momentum transport by the large eddies is inhibited. During half a period (or perhaps one should consider $P/2\\pi$), fluid in such an eddy travels no farther than $(P/2\\tau_l)l$. Zahn (1966, 1989) and \\cite{ZB:89} therefore suppose that the eddy viscosity is reduced by the same factor; that is, they take \\begin{equation} \\nu_\\lambda = \\frac{1}{3}\\lambda v_\\lambda \\min\\left(P/2\\tau_\\lambda,1\\right). \\label{eq:Zlim} \\end{equation} Under this hypothesis, the largest eddies continue to dominate the total viscosity, \\begin{equation} \\nu_{\\rm T}= \\frac{1}{3} v_l l \\min\\left[\\left(\\frac{P}{2\\tau_l}\\right),1\\right]. \\label{eq:Zred} \\end{equation} \\cite{GN:77} and \\cite{GK:77} have taken a different view. These authors argue that the viscosity $\\nu_\\lambda$ on any given scale should be severely suppressed---perhaps by an exponential factor---when $P\\ll 2\\pi\\tau_\\lambda$, so that $\\nu_{\\rm T}$ is dominated by the largest eddies whose turnover time is less than $P/2\\pi$. Since $\\nu_\\lambda\\propto\\tau_\\lambda^2$ [cf. eq.~(\\ref{eq:Kolmogorov}], it follows that \\begin{equation} \\nu_{\\rm T}= \\frac{1}{3} v_l l \\min\\left[\\left(\\frac{P}{2\\pi\\tau_l}\\right)^2,~1\\right]. \\label{eq:Gred} \\end{equation} GN's brief argument can be quoted in its entirety: \\begin{quote} To appreciate this, it is necessary to recognize that turbulent eddies have lifetimes which are comparable to their turnover times. Thus, even though the large convective eddies move across distances of order $l\\tau_{\\rm T}/\\tau_{\\rm c}$ in a tidal period, they do not exchange momentum with the mean flow on this timescale. The contribution to $\\nu_{\\rm T}$ made by the largest eddies is likely to be smaller than $\\nu_{\\rm T~ max}$ by at least one additional factor of $\\tau_{\\rm T}/\\tau_{\\rm c}$. \\end{quote} Underlying this intuition, perhaps, is an analogy with integrable hamiltonian systems, in which secular absorption occurs only through resonances; if such a system is perturbed rapidly compared to any of its autonomous frequencies, no resonance with the perturbing forces is possible. Convective turbulence is neither integrable nor hamiltonian. \\cite{GM:91} have motivated GN's reduction factor by means of an analogy with the kinetic theory of gases, in which they equate $\\tau_l$ to an intermolecular collision time. Unfortunately, as shown in \\S 4, \\cite{VP:96}'s calibration against giants cannot be used to decide between Eqs.~(\\ref{eq:Zred}) and (\\ref{eq:Gred}). Another argument in favor of GN's prescription can be made by considering the time inverse of absorption. Consider a convecting star in an initially exactly circular, short-period orbit. Convection causes small fluctuations in the density distribution within the star and corresponding departures of the exterior potential from symmetry; the latter modulate the binary orbit, albeit at a very low level (\\cite{P:92}). If $\\tau_l\\ll P$, however, the energy of the orbit is adiabatically invariant, so that no secular transfer of energy from the eddies to the orbit can occur. It would therefore seem to follow that the large eddies should not absorb energy from a slightly eccentric short-period orbit. Arguments from time-reversibility are very powerful for thermal systems, but turbulent convection is not reversible because of the flow of energy from large to small scales. So this argument is not conclusive. What is at stake? \\begin{description} \\item[Tidal circularization.] In solar-type stars, $\\tau_l\\sim 20\\dy$ in the middle of the convection zone. Thus $2\\pi\\tau_l$ is substantially longer than the orbital period of circularized main-sequence binaries. Furthermore, the turnover time increases with the depth of the convection zone both on the main sequence towards lower masses, and in the pre-main-sequence phase towards larger radii. \\cite{ZB:89} therefore conclude that pre-main-sequence circularization would be ineffective if the more severe reduction factor (\\ref{eq:Gred}) were correct. \\item[Stellar pulsation] The fundamental pulsation period in variable stars is $\\sim(R^3/GM)^{1/2}$, which is usually shorter than $\\tau_l$ unless the convection zone is thin. \\cite{Gon:82} has argued that the red edge of the instability strip is more consistent with Zahn's prescription (\\ref{eq:Zred}) than with GN's (\\ref{eq:Gred}). \\item[Solar oscillations] Goldreich and collaborators have constructed a theory for the excitation and damping of the solar p-modes by convection [\\cite{GK:77}; \\cite{GK:88}; \\cite{GKM:94} and references therein]. This theory is generally in good agreement with the observed energies of the modes, certainly more so than any alternative, but it would have to be severely modified if eq.~(\\ref{eq:Zred}) were applicable. Acoustic emission by turbulence can be estimated, at least roughly, with some confidence; if the damping processes were enhanced then the energy in the convection would be lower than observed. \\end{description} The outline of our paper is as follows. In the next section we attack the fluid equations directly, deriving linearized equations for the perturbation in the turbulent velocities due to the tide. The in-phase correlation of the perturbed velocity with the tidal force determines the rate of tidal work done on the turbulence. In \\S 3 we study a toy model for tidally perturbed convection: a system of strongly coupled nonlinear oscillators driven by a periodic external force. By numerical integrations, we find the rate of absorption of energy of this system as a function of the frequency of the external forcing. In \\S 4 we discuss the implications of our results for circularization of late-type spectroscopic binaries. ", + "conclusions": "" + }, + "9701/astro-ph9701140_arXiv.txt": { + "abstract": "We consider the effects of advection and radial gradients of pressure and radial drift velocity on the structure of optically thick accretion disks. We concentrate our efforts on highly viscous disks, $\\alpha=1.0$, with large accretion rates. Contrary to disk models neglecting advection, we find that continuous solutions extending from the outer disk regions to the inner edge exist for all accretion rates we have considered. We show that the sonic point moves outward with increasing accretion rate, and that in the innermost disk region advection acts as a heating process that may even dominate over dissipative heating. Despite the importance of advection on it's structure, the disk remains geometrically thin. ", + "introduction": "\\noindent The ``standard accretion disk model'' of Shakura (1972) and Shakura \\& Sunyaev (1973), that has been widely used to model accretion flows around black holes, is based on a number of simplifying assumptions. In particular, the flow is assumed to be geometrically thin and with a Keplerian angular velocity distribution. This assumption allows gradient terms in the differential equations describing the flow to be neglected, reducing them to a set of algebraic equations, and thereby fixes the angular momentum distribution of the flow. For low accretion rates, $\\dot M$, this assumption is generally considered to be reasonable. Since the end of the seventies, however, it has been realized that for high accretion rates, advection of energy with the flow can crucially modify the properties of the innermost parts of accretion disks around black holes. A deviation from a Keplerian rotation may result. Initial attempts to solve the more general disk problem only included advection of energy and the radial gradient of pressure in models with small values of the viscosity parameter, $\\alpha=10^{-3}$ (Paczy\\'nski \\& Bisnovatyi-Kogan 1981), and it was shown that including radial velocity in the radial momentum equation would not change principally the results for such a small $\\alpha$ (Muchotrzeb \\& Paczy\\'nski 1982). Liang \\& Thomson (1980) emphasised the importance of the transonic nature of the radial drift velocity, and the influence of viscosity on the transonic accretion disk solutions was noted by Muchotrzeb (1983), who claimed that such solutions only existed for viscosity parameters smaller than $\\alpha_{*}\\simeq 0.02-0.05$. Matsumoto et al.\\ (1984), then showed that solutions with $\\alpha>\\alpha_{*}$ do in fact exist, but the nature of the singular point, where the radial velocity equals the sound velocity, is changing from a saddle to a nodal type and the position of this point is shifted substantially outwards in the disk. Matsumoto et al.\\ (1984) also demonstrated the non-uniqueness of the solutions with a nodal type critical point for given Keplerian boundary conditions at the outer boundary of the disk (see also Muchotrzeb-Czerny 1986). Extensive investigation of accretion disk models with advection for a wide range of the disk parameters, $\\dot{M}$ and $\\alpha$, was conducted by Abramowicz et al.\\ (1988), with special emphasis on low $\\alpha$. Misra \\& Melia (1996) considered optically thin two-temperature disk models and treated advection in the framework of the Keplerian disk model, but fixed the proton temperature somewhat arbitrarily at the outer boundary. Chakrabarti (1996) solved the advection problem containing shock waves near the innermost disk region, considering accretion through saddle points. Numerical solutions of accretion disks with advection have been obtained by Chen and Taam (1993) for the optically thick case with $\\alpha=0.1$, and by Chen et al.\\ (1996), for the optically thin case (see also Narayan 1996). A simplified account of advection has recently been attempted, either treating it like an additional algebraic term assuming a constant radial gradient of entropy (Abramowicz et al.\\ 1995; Chen et al.\\ 1995; Chen 1995), or using the condition of self-similarity (Narayan \\& Yi 1994). Over the last few years it has become clear, that neglecting the advective heat transport at high $\\dot{M}$ leads to qualitatively wrong conclusions about the topology of the family of solutions of the disk structure equations (see for example Abramowicz et al.\\ 1995; Chen et al.\\ 1995; Artemova et al.\\ 1996). The disk structure equations without advection give rise two branches of solutions: optically thick and optically thin, which do not intersect if $\\dot M<\\dot M_{b} \\approx (0.6-0.9) \\dot{M}_{\\rm Edd}$ for $\\alpha=1$ and $M_{BH}=10^{8}M_{\\odot}$, where $\\dot{M}_{\\rm Edd}$ is the Eddington accretion rate (Artemova et al.\\ 1996). For larger accretion rates there are no solutions of these equations extending continuously from large to small radii, and with Keplerian boundary conditions at the outer boundary of the disk (see also Liang \\& Wandel 1991; Wandel \\& Liang 1991; Luo \\& Liang 1994). It was argued by Artemova et al.\\ (1996), that for accretion rates larger than $\\dot{M}_b$, advection becomes critically important and would allow solutions extending all the way to the inner disk edge also to exist for $\\dot{M}>\\dot{M}_{b}$. The goal of the present paper is to construct explicitly accretion disk models for high $\\dot{M}$ and large $\\alpha$ taking advective heat transport self-consistently into account. We also include radial gradients of pressure and radial drift velocity and we allow for the non-Keplerian character of the circular velocity. Furthermore, we use the geometrically thin disk approximation because, as will be seen in our solutions, the relative thickness of the disk is substantially less than unity. We show that solutions extending from large radii to the inner edge of the disk can be constructed even for accretion rates considerably larger than $\\dot M_b$. We find that advection is very important in the innermost disk region, although the flow does not deviate strongly from Keplerian down to the region where the radial inflow velocity approaches the local sound speed. In \\S 2 we introduce our model and describe our solution methods, while in \\S 3 we discuss our results. ", + "conclusions": "\\noindent In Table 1 we summarise the parameters of the models for which $r_{D}-r_{N}=0$, according to our computations. For each fixed $\\dot{m}$, the properties of the two (or three) self-consistent solutions are similar and differ only quantitatively. In all cases discussed below do we take $M_{BH}=10^{8}M_{\\odot}$ and $\\alpha=1$. We will now compare the solutions with and without advection. In the ``standard model'', for accretion rates $\\dot{m}<\\dot{m}_{b}=14.315$, there always exist solutions that extend continuously from large to small radii. When $\\dot{m}>\\dot{m}_{b}=14.315$ there are no solutions in a range of radii around $r \\approx 13$, and therefore no continuous solutions extending from large radii to the innermost disk edge (see detailed discussion by Artemova et al.\\ 1996, where however, the Newtonian potential was used, resulting in $\\dot{m}_{b}=9.4$). In Figure 2a we plot the disk surface density $\\Sigma=2\\rho h$, in an optically thick disk as a function of radius, $r$, in a model with $\\dot m=10$. The lowest curve is the solution of the standard model, the upper ones are solutions number 2 and 3 in Table 1. Note that the solutions including advection all terminate at radii considerably greater than $r=6$ (inner edge of the disk in the standard model). In Fig.\\ 2b we plot similarly the solutions for $\\dot m=15$. In the standard model, no solution exists in the region around $r\\approx13$, but when advection is included, the structure of the solutions is completely different. Models 6 and 7 in Table 1 are shown. For $\\dot{m}<13$, including the gradient terms gives rather small corrections to the standard disk model, see Figure 2a. When $\\dot{m}>13$ advection becomes essential and for $\\dot{m}>\\dot{m}_{b}$ it changes the picture qualitatively. When $\\dot{m}>\\dot{m}_{b}$ solutions do exist extending continuously from large radii to the innermost disk region where the solution passes through a \"sonic point\" (compare Figs 2a and 2b, see also Fig 4b below). As mentioned above (see the end of Section 2), we can extend our models only down to the region where the radial velocity becomes equal to the local adiabatic sound velocity. We are unable to calculate the properties of the flow for smaller radii. Only more detailed analysis of this region (using other methods) allows one to determine the smoothness of the flow down to the event horizon of a black hole or verify the presence or absence of shocks in the region. In Figure 3 we plot a family of optically thick solutions for different $\\dot{m}$, clearly demonstrating that the solutions to the complete system of disk structure equations including advection and radial gradients have quite different properties at high $\\dot{m}$ compared to the solutions of the standard disk model. In Figure 4a we plot $Q_{\\rm adv}/Q_+$ as a function of radius for $\\dot m=10$ and $\\dot m=28$, that bracket the cases we have studied. Outside the radius where the entropy gradient is zero (and therefore $Q_{\\rm adv}=0$, recall eq.\\ [\\ref{eq:qadv}]), advection provides an additional cooling, that is however, never substantial in our models. On the other hand, inside that radius, advection acts as a heating process that easily dominates over the dissipation rate that decreases rapidly near the inner edge of the disk (as $f\\rightarrow 0$, see eq.\\ [\\ref{eq:qplus}]). Panels 4b and 4c show the corresponding Mach numbers and $h/r$-ratio, respectively. Note that although the flow becomes transonic in the inner region, the disk can still be considered geometrically thin. In our calculations, the non-uniqueness of solutions at large $\\alpha>\\alpha_{*}$, passing through the critical point (Matsumoto et al.\\ 1984; Muchotrzeb-Czerny 1986), is preserved. It is sill not clear, if this non-uniqueness is a realistic physical fact which explanation may be highly problematic (see for example Kato et al.\\ 1988, where the authors argue that the fact that the transonic point is a nodal type critical point is equivalent to an instability condition), or is a result of restrictive precision of our numerical solutions. Two possible approaches to clarify the situation can be suggested. In the first one, we could obtain an asymptotic solution of the disk equations near the gravitational radius and try to match it with the numerical solution going from the nodal point towards the inside. The second approach could be finding stationary solution by solving equations of non-stationary accretion with the appropriate boundary conditions. Both approaches need substantial numerical work, that we plan to undertake in the future." + }, + "9701/astro-ph9701154_arXiv.txt": { + "abstract": "Statler (1987) demonstrated that self-consistent triaxial models with the perfect density law could be constructed for virtually any choice of axis ratios. His experiments are repeated here using triaxial mass models based on Jaffe's density law, which has a central density that diverges as $r^{-2}$, similar to what is observed in low-luminosity elliptical galaxies. Most of the boxlike orbits are found to be stochastic in these models. Because timescales for chaotic mixing are generally shorter than a galaxy lifetime in triaxial models with strong cusps, and because fully-mixed stochastic orbits have shapes that are poorly suited to reproducing a triaxial figure, only the regular orbits are included when searching for self-consistent solutions. As a result of the restriction to regular orbits, self-consistent solutions are found only for mass models with a modest range of shapes, either nearly oblate, nearly prolate or nearly spherical. This result may explain in part the narrow range of elliptical galaxy properties. ", + "introduction": "The grid of 25 mass models is displayed in Figure 1. The short-to-long axis ratio $c/a$ was assigned a value from the set $(0.4,0.5,0.6,0.7,0.8)$, and the triaxiality parameter $T=(a^2-b^2)/(a^2-c^2)$ was set equal to $0.1$ (nearly oblate), $0.3,0.5,0.7,$ or $0.9$ (nearly prolate). For each mass model, 6840 orbits were integrated for 50 dynamical times $T_D(E)$, defined as the period of the 1:1 resonant orbit in the $x-y$ plane at energy $E$. A 7/8 order Runge-Kutta algorithm with variable step size was used for the integrations. Initial conditions were assigned as in Paper I, from one of two grids of starting points: either on an equipotential surface with zero initial velocity (stationary), or in the $x-z$ plane with $v_x=v_z=0$ ($x-z$). Orbits started in the $x-z$ plane are mostly tube orbits, while those started on an equipotential surface are boxlike, passing close to the center. Orbits with both sets of initial conditions were assigned one of a set of 20 energies, defined as the values of the potential on the $x$-axis of a set of ellipsoidal shells -- with the same axis ratios as the density -- that divide the model into 21 sections of equal mass. The grid of starting points at each energy is described in Paper I; it contained 150 orbits per shell in $x-z$ initial condition space, and 192 orbits per shell in stationary initial condition space. Stochasticity was detected by computing the largest Liapunov exponent $\\sigma_1$ for each orbit, defined as the average (over 50 dynamical times) of the exponential rate of divergence between the orbit and an infinitesimally-nearby orbit. At each energy and for each of the initial condition grids, a histogram of the $\\sigma_1$ values was constructed. These histograms invariably exhibited strong bimodality, with one narrow peak near zero (the regular orbits) and another, more diffuse peak centered around larger values of $\\sigma_1$ (the stochastic orbits). As in Paper I, this scheme might sometimes mistake stochastic orbits for regular orbits, since some stochastic orbits remain ``trapped'' near regular parts of phase space for long periods of time. However, any stochastic orbits that are mis-classified in this way are effectively regular, since they can be counted on to mimic regular orbits for many dynamical times. \\section {Orbit Families} Figure 2 gives the average number of stochastic orbits per shell in both initial-condition spaces for each of the 25 mass models. Typically only a few percent of the orbits from the $x-z$ initial condition space were found to be stochastic, but the fraction of stochastic orbits from the stationary initial condition space was much larger, usually more than one-half. While the numbers in Figure 2 can not be simply translated into phase-space fractions, they suggest that stochasticity of the boxlike orbits is most important in strongly triaxial or nearly prolate mass models, and least important in nearly oblate models. (Of course, the boxlike orbits occupy only a small fraction of the total phase space in nearly-axisymmetric models.) Many workers have noted the ``triaxial'' nature of the orbit populations in nearly-prolate models. Figures 3 and 4 illustrate the two initial condition spaces at shell 10 for 12 of the mass models. Regular orbits were assigned to one of three families: long-axis tubes, short-axis tubes and boxes. Long-axis tubes have a nonzero averaged angular momentum about the $x$-axis, short-axis tubes have a definite $L_z$, and boxes have no obvious, time-averaged angular momentum components. The starting points of the three most important resonant orbit families are also plotted in Figures 3 and 4. The 1:1 orbits in the $x-y$ plane generate most of the short-axis tubes; the 1:1 orbits in the $y-z$ plane generate the long-axis tubes; and the $1:2$ $x-z$ ``banana'' orbits generate many of the regular boxlike orbits. As found by Schwarzschild (1993), the number of minor resonances that generate regular orbits in the stationary initial condition space increases as the model becomes rounder. ", + "conclusions": "Strong triaxiality is inconsistent with a high central density. Dynamical models with Jaffe's density law, which has $\\rho\\propto r^{-2}$ near the center, must be nearly oblate ($T=(a^2-b^2)/(a^2-c^2) \\lap 0.4$), nearly prolate ($T\\gap 0.9$) or nearly spherical ($c/a\\gap 0.8$). This result may explain in part the homogeneity of elliptical galaxies, including the tendency of low-luminosity ellipticals to be oblate, and the narrow range of elliptical galaxy kinematical properties. \\bigskip This work was supported by NSF grant AST 90-16515 and by NASA grant NAG 5-2803. G. Quinlan's assistance in the programming is gratefully acknowledged. T. Fridman and M. Valluri kindly read a first draft and made useful suggestions for changes. \\clearpage" + }, + "9701/astro-ph9701224_arXiv.txt": { + "abstract": "Evidence for clustering within the Coma cluster is found by means of a multiscale analysis of the combined angular-redshift distribution. We have compiled a catalogue of 798 galaxy redshifts from published surveys from the region of the Coma cluster. We examine the presence of substructure and of voids at different scales ranging from $\\sim 1$ to $\\sim 16 \\, h^{-1}$ Mpc, using subsamples of the catalogue, ranging from $cz=3000$ km/s to $cz=28000$ km/s. Our substructure detection method is based on the wavelet transform and on the segmentation analysis. The wavelet transform allows us to find out structures at different scales and the segmentation method allows us a quantitative statistical and morphological analysis of the sample. From the whole catalogue we select a subset of $320$ galaxies, with redshifts between $cz=5858$ km/s and $cz=8168$ km/s that we identify as belonging to the central region of Coma and on which we have performed a deeper analysis, on scales ranging from $180\\, h^{-1} $ kpc to $1.44 \\, h^{-1}$ Mpc. Our results are expressed in terms of the number of structures or voids and their sphericity for different values of the threshold detection and at all the scales investigated. According to our analysis, there is strong evidence for multiple hierarchical substructure, on scales ranging from a few hundreds of kpc to about $4 \\, h^{-1} $ Mpc. The morphology of these substructures is rather spherical. On the scale of $720\\, h^{-1} $ kpc we find two main subclusters which where also found before, but our wavelet analysis shows even more substructures, whose redshift position is approximatively marked by these bright galaxies: NGC 4934 \\& 4840, 4889, 4898 \\& 4864, 4874 \\& 4839, 4927, 4875. ", + "introduction": "The Coma cluster (number 1656 in the Abell [1958] catalogue) has been perhaps the most studied galaxy cluster since $1933$, when Zwicky calculated its mass (Zwicky, 1933). It has been long quoted as the paradigmatic example of a roughly spherical, relaxed cluster (Sarazin, 1986). Previous papers (e.g. Fitchett \\& Webster, 1987; Mellier et al., 1988; Baier et al., 1990; Briel et al., 1992; White et al., 1993; Colless \\& Dunn, 1996, hereafter CD96 and Biviano et al., 1996) have suggested that this cluster may have a complex structure. The X-ray images obtained with {\\sl ROSAT} suggest the presence of clumps of emission associated with substructures (Briel et al., 1992; White et al., 1993). However, previous analysis were performed only on 2-D slices of the cluster. In this paper we take up the issue of substructure in this cluster by yet another point of view, namely by trying to make use of redshift information. Our aim is to find out substructure or voids at different scales through a three-dimensional analysis of the cluster, identify them and make a morphological analysis. As it will be evident, our wavelet analysis detects substructure which is not visible in 2-D images, either optical and/or X-ray. The plan of the paper is as follows: in \\S 2 we discuss which selection criteria we have adopted to assemble our catalogue. In \\S 3 and in \\S 4 we discuss the method of analysis based on wavelet transform and segmentation. In \\S 5 we report our results concerning the number and morphology of substructures and in \\S 6 we do the same for the central region of the cluster. In \\S 7 we make some cautionary remarks concerning the statistical and physical significance of our analysis, and finally in \\S 8 we report our conclusions. ", + "conclusions": "During the last years new redhift surveys and methods of analysis have allowed a more through understanding of structure of the Coma cluster (see e.g. Mellier et al., 1988; Escalera et al., 1992; CD96; Biviano et al., 1996), with most of the effort going to ascertain whether it can be classified as a relaxed one or not and to unveil hidden substructures. While this cluster has often in the past been modelled under the assumptions of homogenous velocity structure and spherical symmetry (see e.g. Kent \\& Gunn, 1982), the most recent observational evidence is pointing toward a more complex structure. The recent {\\sl ROSAT} images and 2-D optical analysis have strengthened the evidence for the existence of multiple substructure and suggest that Coma can not be considered a relaxed cluster. Under this respect, it is worth mentioning that already in 1988 Mellier et al. (1988), by analysing the isopleths within a 2-D map of the cluster had suggested the possible existence of $9$ density peaks. In this paper, we have investigated the nature of the Coma cluster performing a 3-D analysis of the combined angular-redshift distribution of the cluster. We have assembled a catalogue of 798 galaxy redshifts, the largest presently available for the Coma cluster. Then, we have developed a 3-D wavelet and segmentation structure analysis that has allowed us to find out substructures on different scales and to describe them in a quantitative way. This powerful method of analysis has already provided excellent results in many fields of physics (e.g. Arneodo et al.,1988; Argoul et al., 1989; Slezak et al., 1990; Fujiwara \\& Soda, 1995; Grebenev et al., 1995). Our results suggests that Coma can not be considered a regular cluster of galaxies, but it is filled up with substructure on all scales ranging from $720$ kpc to $\\sim 4 \\, h^{-1}$ Mpc. The general diminution of the mean morphological parameter, meaning more elongated shapes, and of the number of structures with the scale indicates a hierarchical distribution of the substructure. We have examinated the Coma cluster using three different subsamples of our catalogue (see Table 1); so we have insights within regions of different sizes with different resolutions. On a scale of about $2 \\, h^{-1} $ Mpc, our analysis on the extended Coma catalogue suggests the presence of multiple substructure with spherical morphology (see Table 2). On this scale a large number of voids is detected and their shapes are rather spherical. On the same catalogue, multiple substructure is still present at the scale $4 \\, h^{-1} $ Mpc: on this scale shapes are more elongated ($\\langle L \\rangle =0.67$). On scales $8 \\, h^{-1} $ Mpc and $16 \\, h^{-1} $ Mpc we find only two very elongated objects, in agreement with the histogram of Fig.1a. Voids on scales larger than $2 \\, h^{-1} $ Mpc are few and very elongated. Although we cannot draw any conclusion before having made a comparison with N-body simulations (Pagliaro et al., 1997b), the presence of hierarchically organized substructure seems to point to an evolutionary scenario in which the Coma cluster and the galaxies included in the second peak of the histogram of the galaxy distribution in Fig. 1a were not generated by the collapse of two large spherical density perturbation with different masses and radius of about $15 \\, h^{-1}$ Mpc, but by the merging of a large number of isolated spherical density perturbations of radius ranging from $1 \\, h^{-1}$ Mpc to $3 \\, h^{-1}$ Mpc. This first rough picture of the Coma evolution becomes more evident if we examine the catalogues $C$ and $C_{ext}$. Our analysis of the second catalogue ($C$) suggests the presence of substructures on all the scales with shapes becoming more elongated with growing scale (see Table 3) Voids are detected only on scales $0.47 {\\scriptstyle \\div} 0.94 \\, h^{-1} $ Mpc and their shapes are rather spherical ($0.60 \\le \\langle L \\rangle \\le 0.94$) On smaller scales (hundreds of kpc) we have concentrated our analysis on a central region with redshifts $5858 \\le cz \\le 8168 $ km/s. This region includes the core of Coma with the galaxies NGC 4874 ($\\langle cz \\rangle =$ 7131 km/s) and NGC 4839 ($\\langle cz \\rangle =$ 7397 km/s). Multiple substructures are found on scales $180 {\\scriptstyle \\div} 720 \\, h^{-1} $ kpc with rather spherical morphology ($0.50 \\le \\langle L \\rangle \\le 1.00$). A large number of spherical ($\\langle L \\rangle = 0.98$) voids is detected only on the smallest scale ($180 \\, h^{-1} $ kpc). We stress however once again the fact that the interpretation of substrctures on such small scales in terms of {\\em real} substructures in velocity (or position) space is not as strong, as noted in the previous paragraph. Finally, we concentrate on the scale of $720 \\, h^{-1} $ kpc, where we have detected seven substructures that we show in Fig. 5 and which coincide with the peaks that we find in the central region of the histogram of Fig.1b. To each one of these object we can associate a dominant galaxy. Mean redshifts of the objects are: $cz \\sim 5912$ km/s, $ cz \\sim 6100$ km/s, $cz \\sim 6421$ km/s, $ cz \\sim 6775$ km/s, $cz \\sim 7161$ km/s, $cz \\sim 7594$ km/s and $cz \\sim 7805 $ km/s. In Table 5 we report some statistics only for the clumps containing a significant number of objects. To each clump singled out we can associate one or two dominant galaxies. These are (for increasing redshift): NGC 4934 \\& NGC 4840, NGC 4889, NGC 4898 \\& NGC 4864, NGC 4874 \\& NGC 4839, NGC 4927, NGC 4875. We then confirm the presence of the two subclusters already described by CD96, but in addition we have found statistical evidence for the exstence of more substructures in redshift space. All this evidence leads us to suggest that the Coma cluster can not be considered a regular cluster of galaxies and that its process of formation occurs through a {\\sc bottom-up} mechanism as predicted by {\\sl CDM} and {\\sl MDM} models. Finally, we would like to stress the fact that from this analysis it is difficult to draw any conclusion about the evolutionary state of this cluster. Such an analysis would require some modelling of the evolutionary scenarios through comparison with high-resolution N-body simulations and a better understanding of the velocity field around Coma. We will report on these issues in a subsequent work (Pagliaro et al., 1997b). \\newpage \\begin{table} \\centering \\begin{tabular}{c c c c c c c } \\multicolumn{7}{l}{Structures at 720 $\\, h^{-1}$ kpc} \\\\ \\hline N. Clump & N. Gal. & $cz_{min}$ & $cz_{max}$ & $ $ & $\\sigma_{cz}$ & $M_{vir}$ \\\\ \\hline 1 & 18 & 7756 & 7866 & 7805 & 30 & 6.9 \\\\ 2 & 34 & 7510 & 7666 & 7594 & 39 & 16.8 \\\\ 3 & 71 & 6980 & 7350 & 7161 & 105 & 284 \\\\ 4 & 68 & 6572 & 6967 & 6775 & 102 & 286 \\\\ 5 & 22 & 6356 & 6491 & 6421 & 34 & 10.9 \\\\ 6 & 19 & 6028 & 6164 & 6100 & 36 & 12.4 \\\\ \\hline \\end{tabular} \\caption[h] {Substructures detected on the scale 720 $\\, h^{-1}$ kpc. Redshifts are in km/s. First column gives the number of the clump, the second the number of galaxies contained within the clump, the third the NGC number of the brightest galaxy, and the fourth the values of the masses are in unit of $10^{11} M_{\\odot}$.} \\end{table} \\newpage \\begin{table} \\centering \\begin{tabular}{c c c c c } \\multicolumn{5}{c}{Statistical Tests on $C_{cen}$} \\\\ \\hline scale & $\\langle n \\rangle$ & $\\sigma_{n}$ & ${\\rm n}_{min}$ & ${\\rm n}_{max}$ \\\\ \\hline 0.18 & 23.29 & 1.9 & 17 & 25\\\\ 0.36 & 11.12 & 1.57 & 9 & 15\\\\ 0.72 & 5.05 & 0.72 & 4 & 6\\\\ 1.44 & 1.05 & 0.23 & 1 & 2\\\\ \\hline \\end{tabular} \\caption[h] {Statistics on a number of ``reshufflings'' of redshifts in catalogue $C_{cen}$. In column 1 the scale in $\\, h^{-1}$ kpc is reported, columns 2-5 give the average, standard deviation, minimum and maximum number of structures found, respectively.} \\end{table} \\medskip \\begin{table} \\centering \\begin{tabular}{c c c c c } \\multicolumn{5}{c}{Statistical Tests on $C$} \\\\ \\hline scale & $\\langle n \\rangle$ & $\\sigma_{n}$ & ${\\rm n}_{min}$ & ${\\rm n}_{max}$ \\\\ \\hline 0.47 & 42.12 & 3.79 & 33 & 48\\\\ 0.94 & 17.47 & 3.47 & 12 & 22\\\\ 1.88 & 2.29 & 0.57 & 1 & 3\\\\ 3.76 & 1 & 0 & 1 & 1\\\\ \\hline \\end{tabular} \\caption[h] {Same as Table 6 for catalogue $C$.} \\end{table}" + }, + "9701/astro-ph9701012_arXiv.txt": { + "abstract": "We derive lower bounds on the cosmic baryon density from the requirement that the high-redshift intergalactic medium (IGM) contain enough neutral hydrogen to produce the observed \\lya\\ absorption in quasar spectra. The key theoretical assumption that leads to these analytic bounds is that absorbing structures are no more extended in redshift space than they are in real space. This assumption might not hold if \\lya\\ clouds are highly overdense and thermally broadened, but it is likely to hold in the gravitational instability picture for the \\lya\\ forest suggested by cosmological simulations, independently of the details of the cosmological model. The other ingredients that enter these bounds are an estimate of (or lower limit to) the intensity of the photoionizing UV background from quasars, a temperature $T \\sim 10^4\\;$K for the ``warm'' photoionized IGM that produces most of the \\lya\\ absorption, a value of the Hubble constant, and observational estimates of the mean \\lya\\ flux decrement $\\Dbar$ or, for a more restrictive bound, the distribution function $P(\\tau)$ of \\lya\\ optical depths. With plausible estimates of the quasar UV background and $\\Dbar$, the mean decrement bound implies a baryon density parameter $\\Omega_b \\ga 0.0125 h^{-2}$, where $h \\equiv H_0/100\\;\\kmsmpc$. With conservative values of the UV background intensity and $\\Dbar$, the bound weakens to $\\Omega_b \\ga 0.005 h^{-2}$, but the clustering of the absorbing gas that is required in order to reproduce the observed mean decrement with this baryon fraction is incompatible with other properties of quasar absorption spectra. A recent observational determination of $P(\\tau)$ implies $\\Omega_b \\ga 0.0125 h^{-2}$ even for a conservative estimate of the quasar UV background, and $\\Omega_b \\ga 0.018 h^{-2}$ for a more reasonable estimate. These bounds are consistent with recent {\\it low} estimates of the primordial deuterium-to-hydrogen ratio $\\dtoh$, which imply $\\Omega_b \\approx 0.025 h^{-2}$ when combined with standard big bang nucleosynthesis. Since the bounds account only for baryons in the warm IGM, our results support earlier claims that this component is the dominant reservoir of baryons in the high-redshift universe. The $P(\\tau)$ bound on $\\Omega_b$ is incompatible with some recent high estimates of $\\dtoh$ unless one drops the assumptions of standard big bang nucleosynthesis or abandons the gravitational instability picture for the origin of the \\lya\\ forest. ", + "introduction": "Following the discovery of the first $z>2$ quasar (\\cite{schmidt65}), Gunn \\& Peterson (1965) derived a stringent upper bound on the density of uniformly distributed, neutral hydrogen in intergalactic space, by showing that the redshifted \\lya\\ absorption of neutral gas with more than $\\sim 10^{-4}$ of closure density would turn quasar spectra virtually black at short wavelengths, contrary to observation. They concluded that the intergalactic medium (IGM) must be highly ionized or extremely rarefied. Within a few years, it became clear that the ubiquitous absorption lines in quasar spectra are predominantly those of intervening neutral hydrogen (\\cite{lynds71}; \\cite{sargent80}), and subsequent studies have shown that these lines significantly depress the mean flux received from high-redshift quasars blue-ward of the \\lya\\ emission line (\\cite{oke82}; \\cite{steidel87}; \\cite{jenkins91}; Press, Rybicki, \\& Schneider 1993, hereafter \\cite{press93}; \\cite{zuo93}; \\cite{dobrzycki96}; \\cite{rauch97}). Furthermore, it is now recognized that the ambient ultraviolet (UV) radiation background produced by high redshift quasars will strongly photoionize gas near the cosmic mean density, so that a small amount of diffuse neutral hydrogen corresponds to a much larger amount of total hydrogen (e.g., Haardt \\& Madau 1996, hereafter \\cite{haardt96}). In this paper, we will argue that matching the observed \\lya\\ absorption leads to interesting {\\it lower} bounds on the mean baryon density of the universe, which can be derived from quite general assumptions about the state of the absorbing gas. Recent cosmological simulations suggest that ``\\lya\\ forest'' lines arise in diffuse, but non-uniform, intergalactic gas, and that they therefore represent a phenomenon closely akin to the ``Gunn-Peterson effect'' (\\cite{cen94}; \\cite{petitjean95}; \\cite{zhang95}; \\cite{hkwm96}; for related semi-analytic modeling see \\cite{bi93}; \\cite{bi97}; \\cite{hui97}). Quantitative analyses show that these simulations require a high baryon density in order to reproduce the observed mean opacity of the forest (\\cite{hkwm96}; \\cite{miralda96}; \\cite{cwkh97}; \\cite{rauch97}; \\cite{zhang97}). For the UV background predicted by HM based on the observed population of quasars, matching the mean opacity estimates of PRS typically requires $B \\ga 2$, where \\be B \\equiv \\frac{\\Omega_b h^2}{0.0125} = \\frac{\\eta}{3.4 \\times 10^{-10}} \\label{eqn:Bdef} \\ee is the baryon density scaled to the fiducial big bang nucleosynthesis estimate of Walker et al.\\ (1991). Here $h \\equiv H_0/100\\;\\kmsmpc$, and $\\eta$ is the baryon-to-photon ratio. The bounds on $B$ derived in this paper will not be as high as those derived from the simulations, but they have broader applicability because they are not tied to a specific cosmological scenario, and the simplicity of the arguments that lead to them makes it easier to see how changes in the theoretical or observational inputs affect the final result. We do appeal to the simulations to motivate our one crucial assumption, that structures with a volume filling factor $f$ in real space have, on average, a filling factor no larger than $f$ in redshift space. This assumption can also be phrased as a requirement that typical \\lya\\ forest absorbers satisfy $X \\geq 1$, where \\be X \\equiv \\frac{{\\rm real~space~extent}}{{\\rm redshift~space~extent}} = \\frac{H(z) d}{\\Delta v} \\label{eqn:xdef} \\ee is the ratio of the Hubble flow across an absorber (with line-of-sight extent $d$) to its line width $\\Delta v$. This assumption would not hold in a model of spatially compact \\lya\\ clouds whose line widths are determined by thermal broadening. However, in the physical picture that emerges from the simulations, the marginally saturated ($N_{\\rm HI} \\sim 10^{13}-10^{15}\\;\\cm^{-2}$) absorption lines that dominate the overall absorption usually arise in moderate overdensity structures that are expanding in proper coordinates but contracting in comoving coordinates (\\cite{miralda96}, 1997). The absorption line widths are determined largely by these coherent internal motions rather than by thermal motions, and the velocity extent of these features is generally no larger than the Hubble flow across them. Empirical support for the assumption that $X \\geq 1$ comes from observations of quasar pairs, which yield a typical transverse coherence scale $\\sim 150\\; h^{-1}\\;{\\rm kpc}$ for \\lya\\ forest systems at $z \\sim 2$ (\\cite{bechtold94p}; \\cite{dinshaw94}, 1995). For a typical line width $\\Delta v \\sim 25\\;\\kms$ (\\cite{hu95}), this lengthscale would imply $X \\sim 3/R$ (for $\\Omega=1$, $z=2$), where $R$ is the ratio of transverse extent to line-of-sight extent in real space. If the absorbers are non-spherical, then they are more often intercepted when they are closer to ``face-on,'' so $R$ might reasonably exceed one on average. However, the absorbers would have to be highly flattened, coherent sheets in order to reproduce both the observed transverse coherence and the observed line widths while having an average $X$ significantly smaller than one. The structures giving rise to the \\lya\\ forest in the simulations are commonly filamentary, with the transverse coherence scale corresponding to the thickness of the filaments. Rauch \\& Haehnelt (1995) used the large transverse coherence scale to argue that the \\lya\\ forest must contain a substantial fraction of all baryons in the universe at high redshift. Our arguments here are different from those of Rauch \\& Haehnelt -- in particular, we work directly from observed optical depths instead of from a derived HI column density distribution --- but the spirit is similar. In \\S 2, we consider the lower bound on $B$ that can be obtained from the mean \\lya\\ flux decrement (\\cite{oke82}) alone. In \\S 3 we derive a more restrictive lower bound from the {\\it distribution} of flux decrements (or equivalent optical depths), recently measured from a set of seven Keck HIRES spectra by Rauch et al.\\ (1997). Our analytic approach complements the direct comparison to simulations carried out by Rauch et al.\\ (1997), which leads to stronger but less general bounds on the baryon density. We discuss implications of our results in \\S 4. ", + "conclusions": "The bound on the baryon density from the mean flux decrement, equation~(\\ref{eqn:bound}), and the bound from the optical depth distribution, equation~(\\ref{eqn:bound2}), are the principal results of this paper. The mean decrement bound leads to $B \\ga 1$ if one takes the most convincing estimates of $\\Dbar$ and of the photoionization rate $\\gam$ from quasars. With the Rauch et al. (1997) determination of $P(\\tau)$, the optical depth distribution bound gives $B \\ga 1$ even if one reduces $\\gam$ to half of the best estimate value. These constraints imply a baryon-to-photon ratio $\\eta \\ga 3.4\\times 10^{-10}$ and, assuming the standard model for big bang nucleosynthesis, a corresponding primordial deuterium-to-hydrogen ratio $\\dtoh \\la 6 \\times 10^{-5}$. This constraint is consistent with the estimate $\\dtoh = 2.3 \\pm 0.4 \\times 10^{-5}$ of Tytler, Fan, \\& Burles (1996) from a high-redshift Lyman limit system, but it is inconsistent with the estimates of $\\dtoh \\ga 10^{-4}$ obtained from other high-redshift Lyman limit systems by Songaila et al.\\ (1994), Carswell et al.\\ (1994), and Rugers \\& Hogan (1996ab). Current observational estimates of $P(\\tau)$, $\\gam$, and $h$ can only be reconciled with the high $\\dtoh$ estimates by abandoning the gravitational instability picture of the \\lya\\ forest and returning to a scenario of dense, thermally broadened clouds with $X_w < 1$, or, more radically, by abandoning standard big bang nucleosynthesis. Given an estimated value of $\\gam$, equation~(\\ref{eqn:bound2}) provides an estimate of the number of baryons in the warm ($T \\sim 10^3 - 10^5\\;$K), diffuse ($\\rho/{\\overline \\rho} \\la 10$) intergalactic medium. With the Rauch et al.\\ (1997) $P(\\tau)$, HM's value $\\gam=1.4$, and plausible choices for other parameters (see Fig.~1 caption), one obtains $B_{\\rm IGM} \\approx 1.4$ at $z \\sim 2-3$ (filled circles in Fig.~1). Within standard big bang nucleosynthesis, one must generously increase the estimated observational errors in the primordial $^4{\\rm He}$ abundance even to accommodate $B$ as large as 2 (see, e.g., \\cite{hata95}), and measurements of $({\\rm D}/{\\rm H})$ in the local interstellar medium (see, e.g., \\cite{mccullough92}; \\cite{linsky95}) imply $B \\la 2.4$. This high value of $B_{\\rm IGM}$ therefore suggests that the warm IGM contains most of the baryons in the universe at these redshifts, as predicted by the cosmological simulations and as argued by Rauch \\& Haehnelt (1995) on different, but physically related, empirical grounds. While most of the high-redshift hydrogen resides in the warm, photoionized IGM, most of the {\\it neutral} hydrogen resides in high column density, damped \\lya\\ systems. Because we truncate $P(\\tau)$ at $\\tau_{\\rm max}=3$, equation~(\\ref{eqn:bound2}) severely undercounts the gas in these systems, which contribute $B_{\\rm DLA} \\sim 0.1-0.3 h^2$ at $z \\sim 3$ (\\cite{wolfe95}). Our lower limit to $B_{\\rm IGM}$ far exceeds the baryon density of stars in bright galaxies today, $B \\sim 0.16 h^2$ (\\cite{persic92}). Thus, the baryons that were in the warm IGM at $z=2$ must either (a) remain in the IGM today, (b) have formed brown dwarfs or another form of baryonic dark matter since $z=2$, or (c) have formed stars in systems of very low surface brightness that have been missed in standard estimates of the galaxy luminosity function. There are a number of anticipated observational developments that might strengthen (or weaken) one's confidence in the bounds plotted in Figure~1. The most important will be determinations of $\\Dbar$ and $P(\\tau)$ from larger samples of high resolution, high signal-to-noise ratio quasar spectra, since the values of $\\Bmin$ depend primarily on these observational inputs. Ongoing quasar surveys and, in a few years, the Sloan Digital Sky Survey, will yield improved determinations of the quasar luminosity function, which can be combined with the HM formalism to yield more definitive estimates of the quasar contribution to the photoionizing background. Further analyses of quasar pairs and studies of absorption line shapes in high resolution spectra may provide more compelling evidence for extended \\lya\\ forest absorbers that are broadened largely by Hubble flow, as predicted by cosmological simulations. Such observations would support our key theoretical assumption that $X_w \\ga 1$, and they would reinforce the view that the \\lya\\ forest arises in a smoothly fluctuating intergalactic medium that is the dominant reservoir of high redshift baryons." + }, + "9701/astro-ph9701068_arXiv.txt": { + "abstract": "We present a new method for determining time delays among the light curves of various images in a gravitational lens. The method is based on constructing a simple model for the source variation and forming a $\\chi^2$ measure of the agreement of this same variation with all of the lightcurves. While inspired by Press et al.\\ (1992a, b) our approach is different since we do not assume a Gaussian process for the source variation. Our method has a number of desirable properties: first, it yields an approximate reconstruction of the source variation and of other parameters such as relative time delays; second, it easily incorporates different assumptions about the relations among the light curves and about the data measurement errors; finally, it can be applied to more than two light curves by addition of $\\chi^2$. We apply this method to the light curves of the quadruple gravitational lens PG1115+080 measured by Schechter et al.\\ (1997). Unlike Schechter et al.\\ we include correlated measurement errors in the analysis, as well as the possibility that microlensing may cause different images to vary by different factors in flux. We find a value of $25.0^{+3.3}_{-3.8}$ days ($95\\%$ confidence) for the delay between components B and C (close to the $24$ day value of Schechter et al., and so leading to a similar value of the Hubble constant for a given lens model). However, the ratio $t_{AC}/t_{BA}$ of the two intermediate delays is poorly determined at $1.13^{+.18}_{-.17}$ ($68\\%$ confidence), close to the value predicted by lens models ($\\sim 1.4$) unlike the Schechter et al.\\ value ($\\sim 0.7$). The variation ratios of C with respect to A and of A with respect to B are both different from 1, $1.39^{+.16}_{-.20}$ and $.79^{+.10}_{-.12}$ ($95\\%$ confidence), respectively. This is an indication of a microlensing gradient, and this type of microlensing may allow us to conclude that the size of the quasar optical emission region is about $1000$ AU. ", + "introduction": "Even before the discovery of the first gravitational lens 0957+561 (\\cite{0957dis}) it was recognized that measurements of the time delay between images can be used to determine the Hubble constant (\\cite{refsdal64}, 1966). Despite a history of systematic difficulties, recent measurements combined with an analysis of lens models (\\cite{gn}) indicate that a robust measurement of the Hubble constant ($H_0$) with an accuracy comparable to that of more conventional techniques may be within reach. The measurements include a precise optical time delay (Kundi\\'{c} et al.\\ 1997) which is consistent with the latest results from radio monitoring (Haarsma et al.\\ 1997), and a measurement of the velocity dispersion of the lens galaxy (Falco et al.\\ 1997). Since each lens can potentially yield an independent, single step $H_0$ measurement, obtaining values in a number of lenses with a variety of morphologies and constraints on models may help eliminate systematic errors. Each lens requires a well-constrained lens model, lens and source redshifts, and a measurement of the time delays among the images. Observational constraints on the lens model may include precise image positions and flux ratios as well as data on the objects responsible for the lensing. Resolved structure in images as revealed by radio interferometry provides many constraints, since it is essentially the same as observing multiple sources with the same lens. Whether lensing involves a galaxy, group, or cluster, the position, velocity dispersion, and other observational probes of the mass distribution of the lensing objects all yield constraints on lens models. Measuring time delays requires observing time variations in the image intensities with sufficient accuracy and time resolution. As the number of lenses being carefully monitored increases, more time delays are being determined, such as the preliminary measurement in PKS 1830-211 (\\cite{pks}). Flux measurements are not the only possibility, as shown by the promising measurements of variations in polarization fraction in the images of B 0218+357 (\\cite{iau137}). The quadruply imaged quasar PG1115+080 was the second lens to be discovered (\\cite{weymann}). It is radio quiet, but optical Hubble Space Telescope images (\\cite{hst1115}) were recently analyzed by Schechter et al.\\ (1997, hereafter SCH) and by Keeton \\& Kochanek (1997). They found that lens models which include the effect of the lens galaxy and that of the nearby group of galaxies discovered by Young et al.\\ (1981) can fit the image positions well. They however still found great freedom in the $H_0$ values predicted by these lens models for a given time delay. In four-image configurations where the images lie at roughly the same distance from the lens, there is a well known degeneracy between the radial profile of the lens mass and the inferred $H_0$ (\\cite{deg1}; \\cite{deg2}). In this situation the relative image positions do determine the total enclosed mass within the ring of images, but they are not very sensitive to the radial profile of the mass. Changing the radial profile affects the convergence at the images, and this changes the inferred $H_0$ value in a similar way to the effect of a constant mass sheet (\\cite{falco1}; \\cite{narayan}). In PG1115+080 future observations of the lensing galaxy light profile and, ultimately, a direct measurement of its central velocity dispersion may constrain or eliminate this degeneracy. Recently SCH measured light curves for the A1, A2, B and C images of PG1115+080, and used them to determine multiple delays. The bright A1 and A2 images are close together and should have a very small time delay ($\\sim$ a few hours), so they were combined into a single A=A1+A2 curve. SCH used the Press et al.\\ (1992a, b) method and found that C leads A and A leads B, with $t_{BC}=23.7\\pm3.4$ days and $t_{AC}=9.4\\pm3.4$ days which yields a ratio $r_{ABC}\\equiv t_{AC}/t_{BA}$ of $0.7\\pm0.3$. It is useful to express the two independent quantities as $t_{BC}$ and $r_{ABC}$, since $t_{BC}$ can be taken to fix the Hubble constant for a given lens model while $r_{ABC}$ is independent of overall distance, and can be compared directly with the ratio predicted by lens models. The models mentioned above are consistent in predicting $r_{ABC}=1.4$ to within about $0.1$, and SCH noted the $2\\sigma$ discrepancy with their fitted value. In their analysis SCH assumed that the measurement errors in the light curves are uncorrelated, and also that the fractional flux variations are the same for each component. In this paper we present a more detailed analysis of the light curves in PG1115+080. We first present a new method based on $\\chi^2$ minimization, which has many of the advantages of Press et al.\\ (1992a, b), but is simpler and allows for a more conservative assessment of errors in the reconstructed parameters. We then apply this method to PG1115+080, and include correlated measurement errors in the analysis. In addition to relative time delays, we also allow for different factors of variation in flux, which may arise from microlensing of the images. In \\S 2 we present our $\\chi^2$ method, and discuss its distinct features and free parameters. We then discuss the physical meaning of the various parameters that the method can account for and attempt to extract from data. In \\S 3 we apply our method to fitting the PG1115+080 light curves, singly, in pairs, and all together, and discuss the results and implications. Finally in \\S 4 we summarize our results and point out some of the significant returns possible from further monitoring. ", + "conclusions": "We have developed a method for determining time delays among light curves of multiple images of a gravitational lens. The method constructs a simple model for the actual source variation, using interpolation between a number of equally spaced values. It then performs a combined $\\chi^2$ minimization by fitting all of the light curves to this model simultaneously, which is similar to the method of Press et al.\\ (1992a, b). The ability to vary the basic parameters of the model over a large range lends robustness to our method. Most of the parameters are linear and so the $\\chi^2$ minimization is easily done. In addition to the time delays, the other non-linear parameters are relative variation ratios, which account for different fractional variation in flux for different images. We interpret this physically as evidence for differential microlensing, i.e.\\ a different magnification due to microlensing for the varying region from the region giving rise to the mean flux. Applying our method to the light curves of PG1115+080 observed by SCH, we find a value of $25.0^{+3.3}_{-3.8}$ days ($95\\%$ confidence) for the delay between components B and C, and a ratio $t_{AC}/t_{BA}$ for the two smaller delays of $1.13^{+.18}_{-.17}$ ($68\\%$ confidence). Unlike SCH, we include correlated measurement errors as well as the above mentioned variation ratios in the analysis. Our result for $t_{BC}$ agrees with SCH, but the ratio $r_{ABC}$ does not. Our result for $r_{ABC}$ does agree with lens models, but we find that with the present data it cannot be derived accurately enough to help in fitting lens models. For the variation ratios, we find $\\alpha_{AC}=1.39^{+.16}_{-.20}$ and $\\alpha_{BA}=.79 ^{+.10}_{-.12}$ ($95\\%$ confidence), each indicating differential microlensing at a significance of $4-5$ times our estimated $1\\sigma$ uncertainties. If confirmed as the data accumulates, this would imply that the size of the quasar optical emission region is of order 1000 AU, for microlenses of $\\left=0.1M_{\\sun}$. Further data may also allow a determination of the time variation of the two microlensing magnifications." + }, + "9701/astro-ph9701175_arXiv.txt": { + "abstract": "Results from the Arecibo \\hi Strip Survey, an unbiased extragalactic \\hi survey, combined with optical and 21cm follow-up observations, determine the \\hi Mass Function and the cosmological mass density of \\hi at the present epoch. Both are consistent with earlier estimates, computed for the population of optically selected galaxies. This consistency occurs because, although the distribution of optical central surface brightnesses among galaxies is flat, we fail to find a population of galaxies with central surface brightnesses fainter than 24~$B$-$\\rm mag\\, arcsec^{-2}$, even though there is no observational selection against them. ", + "introduction": "There have been speculations that low surface brightness (LSB) galaxies and intergalactic clouds might constitute a substantial portion of the population of nearby extragalactic objects and that they might contain comparable mass to that in normal galaxies. The LSB population would escape detection optically and would not be included in the galaxy luminosity functions that are used to compute the visible baryonic content of the local Universe (Disney 1976, McGaugh 1996). On the other hand, estimates of the \\hi mass function (HIMF) based on published observations (Briggs 1990) have seemed to indicate that there is probably not any substantial population that has been missed. Weinberg et al (1991) and Szomoru et al (1994) have come to the same conclusion. Until recently, there were no galaxy samples that could be used to address this question empirically, since the galaxies were all first identified optically. In this paper we present results from the Arecibo \\hi Strip Survey, an unbiased 21cm survey with adequate sensitivity to detect \\hi of very low surface density. It is of sufficient length (approximately 15 hours of RA) and depth (7400 \\kms) that it should be immune to fluctuations due to the large scale structure. The total sky coverage was $\\sim\\!65$ square degrees. In the main beam, which has a FWHM of 3.2~arcmin, the survey was capable of detecting \\hi masses of $6\\times 10^5 h^{-2} \\msol$ at 7 $h^{-1}$ Mpc and $1.5\\times 10^8 h^{-2} \\msol$ at the full depth of the survey. The details of the Arecibo Strip Survey are described by Sorar (1994) and Briggs (1996). The survey yielded a total of 61 detections, of which about half could be associated with cataloged galaxies listed in the NASA Extragalactic Database (NED). About five detections with galactic latitude $|b|>10^{\\circ}$, where extinction is not a problem, have no obvious counterparts on the Digitized Sky Survey (DSS). The \\hi selected galaxies generally follow the same structures as optical selected galaxies. This is consistent with (1) results from Szomoru et al (1996) who fail to detect large numbers of \\hi selected galaxies in selected void fields and (2) the finding that LSB galaxies and gas-rich dwarfs lie on structures delineated by normal, high surface brightness galaxies (Mo et al 1994). ", + "conclusions": "" + }, + "9701/astro-ph9701033_arXiv.txt": { + "abstract": "We report the results from an ASCA X-ray observation of the powerful Broad Line Radio Galaxy, 3C109. The ASCA spectra confirm our earlier ROSAT detection of intrinsic X-ray absorption associated with the source. The absorbing material obscures a central engine of quasar-like luminosity. The luminosity is variable, having dropped by a factor of two since the ROSAT observations 4 years before. The ASCA data also provide evidence for a broad iron emission line from the source, with an intrinsic FWHM of $\\sim 120,000\\kmps$. Interpreting the line as fluorescent emission from the inner parts of an accretion disk, we can constrain the inclination of the disk to be $> 35$ degree, and the inner radius of the disk to be $< 70$ Schwarzschild radii. Our results support unified schemes for active galaxies, and demonstrate a remarkable similarity between the X-ray properties of this powerful radio source, and those of lower luminosity, Seyfert 1 galaxies. ", + "introduction": "Unified models of radio sources propose that radio galaxies and radio-loud quasars are basically the same population of objects, viewed at different orientations (Orr \\& Browne 1982; Scheuer 1987; Barthel 1989). The nucleus is only directly visible in quasars, the radio axis of which points within $\\sim45$ degree of the line of sight. In the case of radio galaxies the axis is closer to the plane of the Sky and the nucleus is obscured from view by material in the host galaxy, possibly in a toroidal distribution. The powerful Broad Line Radio Galaxy (BLRG) 3C109 appears to be oriented at an intermediate angle. The nucleus is reddened, $E(B-V) \\sim 0.9$, and polarized in the optical waveband (Rudy \\etal 1984; Goodrich \\& Cohen 1992), suggesting that our line of sight passes through the edge of the obscuring material. The dereddened luminosity of the nucleus, $V=-26.2$ (Goodrich \\& Cohen 1992) identifies the source as an intrinsically luminous quasar. Obscuration is also seen at X-ray wavelengths (Allen \\& Fabian 1992). 3C109 was serendipitously observed with the Position Sensitive Proportional Counter (PSPC) on ROSAT in 1991 August. The PSPC spectrum exhibits soft X-ray absorption in excess of that expected from material within our own Galaxy, implying an intrinsic equivalent hydrogen column density at the redshift of the source ($z = 0.3056$; Spinrad \\etal 1985) of $\\sim 5 \\times 10^{21}$\\apc. The intrinsic (unabsorbed) X-ray luminosity of the source ($0.1-2.4$ keV) determined from the PSPC data is $\\sim 5\\times 10^{45}$\\ergps, making it one of the most X-ray luminous objects within $z\\sim 0.5$; only the QSOs 3C273 and E1821+643 have higher X-ray luminosities (and 3C273 may have a significant beamed component to its X-ray emission). We present here the results of an ASCA X-ray observation of 3C109. The ASCA data confirm the results of Allen \\& Fabian (1992) on excess absorption, and allow us to explore further the X-ray properties of this remarkable source. We show that 3C109 has decreased in brightness by about a factor of two since the ROSAT observations, to a flux level comparable with that observed with the Imaging Proportional Counter (IPC) on the {\\it Einstein Observatory} in 1979 (Fabbiano et al 1984). Also, of particular interest is the detection of a strong, broad iron line in the ASCA spectra. This result implies that most of the X-ray emission from 3C109 is unbeamed. Modelling the line as fluorescent Fe K emission from an accretion disk, we are able to constrain both the inclination and inner radius of the disk. The X-ray properties of 3C109 are shown to be remarkably similar to those of many lower-power, Seyfert 1 galaxies. Throughout this paper we assume a value for the Hubble constant of $H_0$=50 \\kmpspMpc and a cosmological deceleration parameter $q_0$=0.5. ", + "conclusions": "The ASCA results on excess X-ray absorption in 3C109 confirm and refine the earlier ROSAT results (Allen \\& Fabian 1992). The ASCA data show (under the assumption that all of the absorbing material lies at zero redshift) that the X-ray spectrum of the source is absorbed by a total column density of $5.30^{+0.42}_{-0.42} \\times 10^{21}$ \\apc~(Model A). This compares to a Galactic column density of $\\sim 3.0 \\times 10^{21}$ \\apc (Allen \\& Fabian 1996). If we instead assume that the excess absorption, over and above the Galactic value, is due to material at the redshift of 3C109, we determine an intrinsic column density of $4.20^{+0.83}_{-0.78} \\times 10^{21}$ \\apc. Note that these results assume solar abundances in the absorbing material (Morrison \\& McCammon 1983). The X-ray absorption measurements are in good agreement with optical results on the polarization and intrinsic reddening of the source. Goodrich \\& Cohen (1992) determine an intrinsic continuum reddening of $E(B-V) \\sim 0.9$, in addition to an assumed Galactic reddening of $E(B-V) = 0.27$. Using the standard (Galactic) relationship between $E(B-V)$ and X-ray column density, $N_{\\rm H}$/$E(B-V) = 5.8 \\times 10^{21}$ atom cm$^{-2}$ mag$^{-1}$ (Bohlin, Savage \\& Drake 1978), the total reddening observed, $E(B-V) \\sim 1.2$, implies a total X-ray column density (Galactic plus intrinsic) of $\\sim 7.0 \\times 10^{21}$ \\apc. This result is similar to the X-ray column density inferred from the ASCA spectra using model B and confirms the presence of significant intrinsic absorption at the source. Note that this result also suggests that the dust-to-gas ratio in 3C109 is similar to that in our own Galaxy. Further constraints on the distribution of the absorbing gas are obtained from the optical emission-line data presented by Goodrich \\& Cohen (1992). In the narrow line region (NLR), the observed Blamer decrement of H$\\alpha$/H$\\beta = 5.8$ implies (for an assumed recombination ratio of 3.2) an $E(B-V)$ value $\\sim 0.48$. Using the relationship of Bohlin, Savage \\& Drake (1978) this implies an X-ray column density to the NLR of $\\sim 2.8 \\times 10^{21}$ \\apc, in good agreement with the Galactic column density of $\\sim 3.0 \\times 10^{21}$ \\apc~determined by Allen \\& Fabian (1996) and adopted in the X-ray analysis presented here. The Balmer decrement in the broad line region (BLR) is very steep (H$\\alpha$/H$\\beta = 13.2$). Although this value cannot be reliably used to infer the extinction to the BLR, the intrinsic line ratio is unlikely to be above 5, suggesting a total line-of sight reddening to the BLR of $E(B-V) \\approxgt 0.8$. Thus, the BLR is likely to be intrinsically reddened by $E(B-V) \\approxgt 0.3$. The optical emission line results are therefore consistent with the two-component absorber model (B), with the column density of the intrinsic absorber being comparable with the Galactic component. 3C109 is the most powerful object in which a strong broad iron line has been resolved to date. Several more luminous quasars observed with ASCA do not show any iron emission or reflection features (Nandra et al 1995). The next most luminous object with a confirmed broad line is 3C390.3 (Eracleous, Halpern \\& Livio 1996) which is about 10 times less luminous in both the X-ray and radio bands than 3C109. The equivalent widths of the lines in both objects are $\\sim 300$ eV and therefore similar to those observed in lower-luminosity Seyferts. This argues against any X-ray `Baldwin effect' (as proposed by Iwasawa \\& Taniguchi 1993). The line emission from 3C109 is most plausibly due to fluorescence from the innermost regions of an accretion disc around a central black hole (Fabian et al 1995). Our results constrain the inner radius of the accretion disk to be $< 70 R_{\\rm s}$ and the inclination of the disk to be $ > 35$ degree. The strong iron line observed in 3C109, and the lack of evidence for a synchrotron self-Compton continuum in the X-ray spectrum, both suggest that little radiation from the jet is beamed into our line of sight. The inclination determined from the ASCA data is larger than the angle proposed by Giovannini \\etal (1994) based on the jet/core flux ratio of the source ($\\theta < 34$ degree). However, the jet/core flux arguments are based on simple assumptions about the average orientation angles for radio galaxies and neglect environmental effects. The conflict with the X-ray results may indicate that the situation is more complicated. Giovannini \\etal (1994) also present constraints on the inclination from VLBI observations of the jet/counterjet ratio, which require $\\theta < 56$ degree. The VLBI constraint, together with the ASCA X-ray constraint, then suggests $35 < \\theta < 56$ degree. Our results on 3C109 are in good agreement with the unification schemes for radio sources and illustrate the power of X-ray observations for examining such models. The preferred, intermediate inclination angle for the disk in 3C109 is in good agreement with the results on X-ray absorption, polarization and optical reddening of the source, all of which suggest that our line of sight to the nucleus passes closes to the edge of the surrounding molecular torus. The results on the broad iron line reveal a striking similarity between the X-ray properties of 3C109 and those of lower power, Seyfert 1 galaxies (Mushotzky et al 1995; Tanaka et al 1995; Iwasawa et al 1996). This is despite the fact that the X-ray power of 3C109 exceeds that of a typical Seyfert galaxy by $\\sim 2$ orders of magnitude." + }, + "9701/astro-ph9701205_arXiv.txt": { + "abstract": "An area large enough (180 $\\mbox{arcmin}^{2}$) to put constraints on a possible low mass brown dwarf population in the Pleiades has been surveyed to very faint magnitudes in $I$, $J$ and $K$. The completeness limit, I=21.6, corresponds to a mass of 0.01 $\\mbox{M}_{\\odot}$ for a cluster age of 70 Myr and 0.035 $\\mbox{M}_{\\odot}$ for 120 Myr. The result is consistent with previous investigations at higher masses that the brown dwarf initial mass function is a $m^{-1}$, or even less steep, power law. Thus low mass brown dwarfs cannot contribute significantly to the Pleiades' mass. One new possible Pleiades member was found, mass $\\sim0.08$ $\\mbox{M}_{\\odot}$ (age 120Myr). ", + "introduction": "Brown dwarfs (BDs) are stellar-like objects. The only difference from ordinary stars is that the mass is too low to bring up the central temperature to the level of stable hydrogen burning, thus the BD luminosity decreases with time. As an example, from 70 Myr to 10 Gyr, a 0.08 $\\mbox{M}_{\\odot}$ object at the hydrogen burning limit would decrease a factor 15 in luminosity, while a 0.06 $\\mbox{M}_{\\odot}$ would go a factor 700 (Burrows et al. 1993). The ideal target for a BD search thus is a fairly young, nearby and rich star cluster. The Pleiades is the obvious choice in the northern hemisphere, being at $\\sim 125$ pc and 70-120 Myr old. Several recent authors have proposed an age above 100 Myr. In this paper 120 Myr is adopted. The mean distance modulus of the cluster used is 5.53 (see Basri et al. 1996 for a discussion). Whether BDs could be a significant part of the local dark matter is a subject of controversy. Several recent photometric surveys have found a significant drop in the luminosity function from $M_{V}=12$ to $M_{V}=14$, leading to a turnover in the initial mass function (IMF, $dN_{\\rm stars} = const*m^{-n}*dm$ ($m$ = mass; $n$ = IMF-index)) at $\\sim$ 0.3 $\\mbox{M}_{\\odot}$ (see e.g. Gould et al. 1996; Tinney 1993) or a continued rise towards the hydrogen burning limit (see e.g. Kirkpatrick et al. 1994) depending on the choice of mass-luminosity relation. Kroupa (1995) showed that the difference between the nearby stellar luminosity function (LF), measured by parallax and the more distant LF, measured by photometry alone, can be explained by undetected binary companions in the distant sample. {}From a recently derived mass-luminosity relation (Chabrier et al. 1996) and the well-known photometric LF (see e.g. Gould et al. 1996), Mera et al. (1996) concluded that the IMF for low mass stars continues to rise to the hydrogen burning limit. However the number of known field stars at the low mass end of the main sequence is small. To clearify this item it is necessary to discover more low mass stars and BDs, preferrably at a known distance and age. In \\S 2, observations, reductions and a short discussion on photometry and completeness limits is given. In \\S 3, the extraction of Pleiades members is described. \\S 4 discusses contamination and overall observing strategy. In \\S 5 the implications of this paper on the IMF-index is discussed and compared to other authors. ", + "conclusions": "Our goal was not to cover as large an area as possible, but to reach very faint magnitudes in order to investigate the presence of a population of low mass BDs in the Pleiades. This survey is complete to $I = 21.6$, corresponding to 0.035 $\\mbox{M}_{\\odot}$ (120 Myr) or 0.01 $\\mbox{M}_{\\odot}$ (70 Myr). In Fig. 3, LFs from recent surveys are compared to LFs deduced from 4 different IMF-indices. \\begin{figure} \\picplace{6.95 cm} \\caption{Comparing LF predicted for four different IMF-indices (bars) and five recent surveys (curves). All LFs were normalized to 10 counts in $1610^3$ & $10$ & $0.01$? \\\\ & & & \\\\ GMCs & $>10^3$ & $15-40$ & \\\\ H~II regions & $>10$ & $10^4$ & \\\\ & & & \\\\ SNR & $>1$ & $10^4-10^7$ & \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{flushleft} \\end{table} The gaseous component of this medium should be reasonably well described as a compressible, self-gravitating fluid (e.g, Shu 1992, Ch. 1) obeying the magnetohydrodynamic (MHD) equations, since even very low ionization fractions allow coupling of the motions to the magnetic field (Mestel \\& Spitzer 1956; see also Shu 1992, Ch.\\ 27). From one of these equations, a very important result called the Virial Theorem (VT) can be obtained, which describes the energy balance (or ``budget'') of particular regions within the medium. In this course, we first review the derivation of the VT (\\S\\ \\ref{derivation}), stressing in particular the physical meaning of each term, including the often-neglected surface terms, and the validity of frequently used simplifying assumptions (e.g., Shu 1992). In \\S\\ \\ref{applications} we then discuss various applications of the Virial Theorem, both to idealized cases and to real interstellar cloud properties, pointing out existing incompletenesses. Finally, in \\S\\ \\ref{conclusions} contains the conclusions. Throughout this course, we shall ignore the flux-loss effects of ambipolar diffusion (e.g, Shu 1992). ", + "conclusions": "From the discussion in the previous sections, a number of points can be concluded. \\medskip \\noindent 1. ``Virial equilibrium'', i.e., $d^2I/dt^2=0$ allows up to a constant rate of change in the moment of inertia $I$, even though it is commonly taken to mean a stationary equilibrium. \\medskip \\noindent 2. Larson's (1981) relations, eqs.\\ (\\ref{larvel}) and (\\ref{lardens}), have been interpreted as a consequence of virial equilibrium\\footnote{Mouschovias \\& Psaltis (1995) distinguish between magnetic virial equilibrium and magnetic support, but the resulting balance equation is the same.} between self-gravity on the one hand and turbulent velocity dispersion and/or magnetic support on the other. From the observational data discussed above, this interpretation may apply to strongly condensed clumps which, due to self-gravity, have decoupled from their surrounding medium, and thus require internal support against collapse. \\medskip \\noindent 3. The above interpretation, however, is insufficient for explaining {\\it both} scaling relations. One is still independent and requires an additional explanation. A critical mass-to-flux ratio has been invoked as the origin of the density-size relation assuming roughly constant magnetic field strengths (Shu et al.\\ 1987; Myers \\& Goodman 1988b; Mouschovias \\& Psaltis 1995), but this assumption is not verified in the observational data. For example, the magnetic field strenghts discussed by Myers \\& Goodman (1988a) and Mouschovias and Pslatis (1995) span three orders of magnitude. Besides, examples of virialized clouds which however follow different scaling relations exist (e.g., Loren 1989; Fuller \\& Myers 1992; Caselli \\& Myers 1995; Goodman et al.\\ 1997; see also V\\'azquez-Semadeni \\& Gazol 1995 for a theoretical discussion). Planar shocks have also been invoked as an explanation of the apparently constant column densities (Larson 1981; Scalo 1987). \\medskip \\noindent 4. It is possible that the density-size relation is not a true property of clouds, but only an observational effect due to the limited integration times used in surveys (Larson 1981; Kegel 1989; Scalo 1990), which tend to select constant-column density obejcts. In this case, surveys using larger integration times would exhibit an increasingly larger scatter in density-size plots, as in the data of Falgarone et al.\\ (1992) or the numerical data of V\\'azquez-Semadeni et al.\\ (1997), the latter being free from such limitations. In fact, large scatter was already present in the Myers \\& Goodman (1988b) data. Although they interpreted it as uncertainty in the data rather than real scatter about virialization, the latter possibility is equally feasible. A study that seemed to find exceptionally constant column densities while claiming a large dynamic range (Wood et al.\\ 1994) has been questioned by \\VS\\ et al.\\ (1997). \\medskip \\noindent 5. The velocity dispersion-size relation does not appear subject to the possibility of a spurious origin, and therefore it is probably real {\\it when detected}, although in many cases it is not (e.g., Loren 1989; Plume et al.\\ 1997). For the cases when it is present, a number of physical mechanisms are plausible candidates. If the density-size relation is not verified, then the standard arguments based on virial equilibrium between gravity and velocity dispersion cannot be invoked (recall the $\\sigma$-$R$ relation is a consequence of virial equilibrium {\\it and} the density-size relation). In particular, the possibility that the $\\sigma$-$R$ relation is a consequence of the characteristic energy spectrum of an ensemble of shocks in the flow has been suggested by a number of authors (e.g., Passot et al.\\ 1988; Padoan 1995; Gammie \\& Ostriker 1996; Fleck 1996; \\VS\\ et al.\\ 1997). In this case, the origin of the dispersion-size relation would be completely independent of the virial condition, explaining why the relation can be present even in the absence of a density-size relation. \\medskip \\noindent 6. In summary, the Larson's relations and their virial equilibrium interpretation probably apply to relaxed, strongly self-gravitating clumps. Regions which do not exhibit clear scaling relations often include either massive cores (Caselli \\& Myers 1995; Plume et al.\\ 1997) or regions with strong evidence of recent perturbations (Loren 1989). In fact, the dense cores studied by Plume et al.\\ were selected by the presence of H$_2$O masers, suggesting strong excitation mechanisms as well. Such ``perturbed'' regions are likely not in virial equilibrium, thus also being transients (e.g., Magnani et al.\\ 1993), or at least having strongly fluctuating moments of inertia (shapes and internal mass distributions). \\vspace*{0.3cm}" + }, + "9701/astro-ph9701080_arXiv.txt": { + "abstract": "The X-ray satellite BeppoSAX \\footnote{ The BeppoSAX team is composed by scientists from: \\begin{itemize} \\item Istituto Astrofisica Spaziale (IAS), C.N.R., Frascati and Unita' GIFCO Roma \\item Istituto di Fisica Cosmica ed Applicazioni Informatica (IFCAI), C.N.R., and Unita' GIFCO, Palermo \\item Istituto Fisica Cosmica e Tecnologie Relative (IFCTR), C.N.R. and Unita' GIFCO, Milano \\item Istituto per le Tecnologie e Studio Radiazioni Extraterrestri (ITeSRE), C.N.R., Bologna and Universita' di Ferrara \\item Space Research Organization of the Netherlands (SRON), The Netherlands \\item Space Science Department (SSD), ESA, Noordwijk, The Netherlands \\item BeppoSAX Science Data Center, Rome \\item BeppoSAX Science Operation Center, Rome \\end{itemize}} , a major programme of the Italian space agency with participation of the Dutch space agency, was successfully launched from Cape Canaveral on April 30, 1996. After a 2 months period devoted to engineering check out that confirmed the nominal functionality of the satellite and the scientific payload, we have performed a series of observations of celestial objects devoted to calibrate the instruments and verify their scientific performances. Here we will present some preliminary results obtained in this phase. They are confirming the expected scientific capabilities of the mission. ", + "introduction": "The X-ray satellite SAX, named BeppoSAX after launch in honour of Giuseppe (Beppo) Occhialini, is the first X-ray mission with a scientific payload covering more than three decades of energy - from 0.1 to 300 keV - with a relatively large area, a good energy resolution, and with imaging capabilities (resolution of about 1 arcmin) in the range of 0.1-10 keV. This capability, in conjunction with the presence of wide field instruments primarily aimed at discovering transient phenomena, which could then be observed with the broad band instruments, provides an unprecedented opportunity to study the broad band behaviour of several classes of X-ray sources. The broad band capability is provided by a set of instruments co-aligned with the Z axis of the satellite, Narrow Field Instruments (hereafter NFI) and composed by: \\begin{itemize} \\item MECS (Medium Energy Concentrator Spectrometers): a medium energy (1.3-10 keV) set of three identical grazing incidence telescopes with double cone geometry (Citterio et al. 1985, Conti et al. 1994), with position sensitive gas scintillation proportional counters in their focal planes (\\cite{mecs}). \\item LECS (Low Energy Concentrator Spectrometer): a low energy (0.1-10 keV) telescope, identical to the other three, but with a thin window position sensitive gas scintillation proportional counter in its focal plane (\\cite{lecs}). \\item HPGSPC, a collimated High Pressure Gas Scintillation Proportional Counter (4-120 keV, \\cite{hpgspc}). \\item PDS, a collimated Phoswich Detector System (15-300 keV, \\cite{pds}) \\end{itemize} Access to large regions of the sky ($\\sim3000$ degree$^2$) with a resolution of 5' in the range 2-30 keV is provided by: \\begin{itemize} \\item two coded mask proportional counters (Wide Field Cameras, WFC, \\cite{wfc}), perpendicular to the axis of the NFI and pointed in opposite directions. \\end{itemize} Finally, the anticoincidence scintillator shields of the PDS (GRBM) will be used as a gamma-ray burst monitor in the range 60-600 keV with a fluence greater than about $10^{-6}$ erg cm$^{-2}$ and with a temporal resolution of about 1 ms. More details on the mission and its instruments can be found in \\cite{psb95}, \\cite{general}, in the special session devoted to BeppoSAX of the SPIE Vol. 2517 and on line at: {\\it http://www.sdc.asi.it.} ", + "conclusions": "" + }, + "9701/hep-ph9701304.txt": { + "abstract": "At the end of the inflationary stage of the early universe, profuse particle production leads to the reheating of the universe. Such explosive particle production is due to parametric amplification of quantum fluctuations for the unbroken symmetry case (appropriate for chaotic inflation), or spinodal instabilities in the broken symmetry phase (which is the case in new inflation). This mechanism is non-perturbative and depends on the details of the particle physics models involved. A consistent study of this mechanism requires a detailed analysis and numerical treatment with an approximation scheme that ensures energy (covariant) conservation and a consistent non-perturbative implementation. We study the $ O(N) $ symmetric vector model with a quartic self-interaction in the large $ N $ limit, Hartree and resummed one-loop approximations (with $ N = 1 $) to address the non-perturbative issues. The non-equilibrium equations of motions, their renormalization and the implementation of the approximations are studied in arbitrary spatially flat FRW cosmologies. A full description, analytically and numerically is provided in Minkowski space-time to illustrate the fundamental phenomena in a simpler setting. We give analytic results for weak couplings and times short compared to the time at which the fluctuations become of the same order as the tree level terms, as well as numerical results including the full backreaction. In the case where the symmetry is unbroken, the analytical results agree spectacularly well with the numerical ones in their common domain of validity. In the broken symmetry case, interesting situations, corresponding to slow roll initial conditions from the unstable minimum at the origin, give rise to a new and unexpected phenomenon: the dynamical relaxation of the vacuum energy. That is, particles are abundantly produced at the expense of the quantum vacuum energy while the zero mode comes back to almost its initial value. We obtain analytically and numerically the equation of state which in both cases can be written in terms of an effective polytropic index that interpolates between vacuum and radiation-like domination. The self-consistent methods presented in these lectures are the only approaches, so far, that lead to reliable quantitative results on the reheating mechanism in the inflationary universe. These approaches take into account the non-linear interaction between the quantum modes and exactly conserve energy (covariantly). Simplified analysis that do not include the full backreaction and do not conserve energy, result in unbound particle production and lead to quantitatively erroneous results. For spontaneously broken theories the issue of whether the symmetry may be restored or not by the quantum fluctuations is analyzed. The precise criterion for symmetry restoration is presented. The field dynamics is symmetric when the energy density in the initial state is larger than the top of the tree level potential. When the initial energy density is below the top of the tree level potential, the symmetry is broken. Finally, we provide estimates of the reheating temperature as well as a discussion of the inconsistency of a kinetic approach to thermalization when a non-perturbatively large number of particles is created. ", + "introduction": "Research activity on inflationary cosmologies has continued steadily since the concept of inflationary cosmology was first proposed in 1981 \\cite{guth}. It was recognized that in order to merge an inflationary scenario with standard Big Bang cosmology a mechanism to reheat the universe was needed. Such a mechanism must be present in any inflationary model to raise the temperature of the Universe at the end of inflation, thus the problem of reheating acquired further importance deserving more careful investigation. The original version of reheating envisaged that during the last stages of inflation when the universe expansion slows down, the energy stored in the oscillations of the inflaton zero mode transforms into particles via single particle decay. Such particle production reheats the universe whose temperature was redshifted to almost zero during the inflationary expansion \\cite{newI}. It was realized recently\\cite{frw,stb,lindekov,jap}, that the elementary theory of reheating \\cite{newI} does not describe accurately the quantum dynamics of the fields when the oscillations of the inflaton field (zero mode) have large amplitude. \\bigskip Our programme on non-equilibrium dynamics in quantum field theory, started in 1992\\cite{boyveg}, is naturally poised to provide a framework to study these problems. The larger goal of the program is to study the dynamics of non-equilibrium processes, such as phase transitions, from a fundamental field-theoretical description, by obtaining and solving the dynamical equations of motion for expectation values and correlation functions of the underlying four dimensional quantum field theory for physically relevant problems: phase transitions and particle production out of equilibrium, symmetry breaking and dissipative processes. The focus of our work is to describe the quantum field dynamics when the energy density is {\\bf large}. That is, a large number of particles per volume $ m^{-3} $, where $ m $ is the typical mass scale in the theory. Usual S-matrix calculations apply in the opposite limit of low energy density and since they only provide information on {\\em in} $\\rightarrow$ {\\em out} matrix elements, are unsuitable for calculations of time dependent expectation values. Our methods were naturally applied to different physical problems like pion condensates \\cite{dcc,linon,inh}, supercooled phase transitions \\cite{boyveg,dis,rev}, inflationary cosmology \\cite{frw,dis,rev,big,desit,pfrw}, the hadronization stage of the quark-gluon plasma\\cite{muller} as well as trying to understand out of equilibrium particle production in strong electromagnetic fields and in heavy ion collisions \\cite{boyveg,dcc,losalam}. When a large energy density is concentrated in one or few modes, for example the inflaton zero mode in inflationary cosmology, under time evolution this energy density will be transferred to other modes driving a large amplification of quantum fluctuations. This, in turn, gives rise to profuse particle production for bosonic fields, creating quanta in a highly non-equilibrium distribution, radically changing the standard picture of reheating the post-inflationary universe\\cite{newI,dolgov}. Fermionic fields are not very efficient for this mechanism of energy ``cascading'' because of Pauli blocking \\cite{linon}. The detail of the processes giving rise to preheating can be different depending on the potential for the scalar field and couplings to other fields involved, as well as the initial conditions. For example, in new inflationary scenarios, where the expectation value of the zero mode of the inflaton field evolves down the flat portion of a potential admitting spontaneous symmetry breaking, particle production occurs due to the existence of unstable field modes whose amplitude is amplified until the zero mode leaves the instability region. These are the instabilities that give rise to spinodal decomposition and phase separation. In contrast, if we start with chaotic initial conditions, so that the field has large initial amplitude, particles are created from the parametric amplification of the quantum fluctuations due to the oscillations of the zero mode and the transfer of energy to higher modes. In these lectures we analyze the details of this so-called {\\bf preheating} process both analytically as well as numerically. Preheating is a non-perturbative process, with typically $ 1\\slash \\lambda $ particles being produced, where $\\lambda$ is the self coupling of the field. Due to this fact, any attempts at analyzing the detailed dynamics of preheating must also be non-perturbative in nature. This leads us to consider the $ O(N) $ vector model in the large $ N $ limit. This is a non-perturbative approximation that has many important features that justify its use: unlike the Hartree or mean-field approximation\\cite{dis}, it can be systematically improved in the $ 1\\slash N $ expansion. It conserves energy, satisfies the Ward identities of the underlying symmetry, and again unlike the Hartree approximation it predicts the correct order of the transition in equilibrium. This approximation has also been used in other non-equilibrium contexts\\cite{boyveg,dcc,losalam}. Our main results can summarized as follows \\cite{big}. We provide consistent non-perturbative analytic estimates of the non-equilibrium processes occurring during the preheating stage taking into account the {\\bf exact} evolution of the inflaton zero mode for large amplitudes when the quantum back-reaction due to the produced particles is negligible i.e. at early and intermediate times. We also compute the momentum distribution of the number of particles created, as well as the effective equation of state during this stage. Explicit expressions for the growth of quantum fluctuations, the preheating time scale, and the effective (time dependent) polytropic index defining the equation of state are given in sec. IV and V. We go beyond the early/intermediate time regime and evolve the equations of motion numerically, taking into account back-reaction effects. (That is, the non-linear quantum field interaction). These results confirm the analytic estimates in their domain of validity and show how, when back-reaction effects are large enough to compete with tree level effects, dissipational effects arise in the zero mode. Energy conservation is guaranteed in the full backreaction problem, leading to the eventual shut-off of particle production. This is an important ingredient in the dynamics that determines the relevant time scales. We also find a novel dynamical relaxation of the vacuum energy in this regime when the theory is in the broken phase. Namely, particles are produced at the expense of the quantum vacuum energy while the zero mode contributes very little. We find a radiation type equation of state for late times ($ p \\approx \\frac13 \\; \\varepsilon $) despite the lack of local thermodynamic equilibrium. Finally, we provide an estimate of the reheating temperature under clearly specified (and physically reasonable assumptions) in a class of models. We comment on when the kinetic approach to thermalization and equilibration is applicable. There have been a number of papers (see refs.\\cite{stb,lindekov,jap} -\\cite{tkachev}) dedicated to the analysis of the preheating process where particle production and back-reaction are estimated in different approximations \\cite{paris}. Our analysis differs from other works in many important aspects. We emphasize the need of a non-perturbative, self-consistent treatment that includes backreaction and guarantees energy conservation (covariant conservation in the expanding universe) and the conservation of all of the important symmetries. Although analytic simplified arguments may provide a qualitative picture of the phenomena involved, a quantitative statement requires a detailed numerical study in a consistent manner. Only a self-consistent, energy conserving scheme that includes backreaction effects can capture the corresponding time scales. Otherwise infinite particle production may result from uncontrolled approximations. The layout of these lectures is as follows. Section II presents the model, the evolution equations, the renormalization of the equations of motion and introduces the relevant definitions of particle number, energy and pressure and the details of their renormalization. The unbroken and broken symmetry cases are presented in detail and the differences in their treatment are clearly explained. In sections III through V we present a detailed analytic and numerical treatment of both the unbroken and broken symmetry phases emphasizing the description of particle production, energy, pressure and the equation of state. In the broken symmetry case, when the inflaton zero mode begins very close to the top of the potential, we find that there is a novel phenomenon of relaxation of the vacuum energy that explicitly accounts for profuse particle production through the spinodal instabilities and energy conservation. We discuss in section VI why the phenomenon of symmetry restoration at preheating, discussed by various authors\\cite{lindekov,tkachev,kolbriotto,kolblinde} is {\\bf not} seen to occur in the cases treated by us in ref.\\cite{dis,big} and relevant for new inflationary scenarios \\cite{paris}. A precise criterion for symmetry restoration is given. The symmetry is broken or unbroken depending on the value of the initial energy density of the state. When the energy density in the initial state is larger than the top of the tree level potential then the symmetry is restored \\cite{paris}. When it is smaller than the top of the tree level potential, then it is broken and Goldstone bosons appear \\cite{dis,big}. In the first case, the amplitude of the zero mode is such that $ V(\\eta_0) > V(0) $ (all energy is initially on the zero mode). In this case the dynamics is very similar to the unbroken symmetry case, the amplitude of the zero mode will damp out, transferring energy to the quantum fluctuations via parametric amplification, but asymptotically oscillating around zero with a fairly large amplitude. In section VII we briefly discuss the amplitude expansion (linearizing in the field amplitude) and compare with a full non-linear treatment in a model for reheating where the inflaton decays into a lighter scalar field \\cite{linon}. In section VIII we provide estimates, under suitably specified assumptions, of the reheating temperature in the $ O(N) $ model as well as other models in which the inflaton couples to lighter scalars. In this section we argue that thermalization cannot be studied with a kinetic approach because of the non-perturbatively large occupation number of long-wavelength modes. Finally, we summarize our results and discuss future avenues of study in the conclusions. ", + "conclusions": "It is clear that preheating is both an extremely important process in a variety of settings, as well as one involving very delicate analysis. In particular, its non-perturbative nature renders any treatment that does not take into account effects such as the quantum back-reaction due to the produced particles, consistent conservation (or covariant conservation) of the relevant quantities and Ward identities, incapable of correctly describing the important physical phenomena during the preheating stage. In this work, we dealt with these issues by using the $O(N)$ vector model in the large $N$ limit. This allows for a controlled non-perturbative approximation scheme that conserves energy and the proper Ward identities, to study the non-equilibrium dynamics of scalar fields. Using this model we were able to perform a full analysis of the evolution of the zero mode as well as of the particle production during this evolution. Our results are rather striking. We were able to provide analytic results for the field evolution as well as the particle production and the equation of state for all these components in the weak coupling regime and for times for which the quantum fluctuations, which account for back-reaction effects, are small. What we found is that, in the unbroken symmetry situation, the field modes satisfy a Lam\\'e equation that corresponds to a Schr\\\"odinger equation with a two-zone potential. There are two allowed and two forbidden bands, which is {\\em decidedly} unlike the Mathieu equation used in previous analysis\\cite{stb,lindekov,jap}. The difference between an equation with two forbidden bands and one with an infinite number is profound. We were also able to estimate analytically the time scale at which preheating would occur by asking when the quantum fluctuations as calculated in the absence of back-reaction would become comparable to the tree level terms in the equations of motion. The equations of state of both the zero mode and ``pions'' were calculated and were found to be describable as polytropes. These results were then confirmed by numerical integration of the equations, and we found that the analytic results were in great agreement with the numerical ones in their common domain of validity. When the $O(N)$ symmetry is spontaneously broken, more subtle effects can arise, again in the weak coupling regime. If the zero mode starts off very near the origin, then the quantum back-reaction grows to be comparable to the tree level terms within one or at most a few oscillations even for very weak coupling. In this case the periodic approximation for the dynamics of the zero mode breaks down very early on and the full dynamics must be studied numerically. When numerical tools are brought to bear on this case we find some extremely interesting behavior. In particular, there are situations in which the zero mode starts near the origin (the initial value depends on the coupling) and then in one oscillation, comes back to almost the same location. However, during the evolution it has produced $1\\slash g$ particles. Given that the total energy is conserved, the puzzle is to find where the energy came from to produce the particles. We found the answer in a term in the energy density that has the interpretation of a ``vacuuum energy'' that becomes {\\em negative} during the evolution of the zero mode and whose contribution to the equation of state is that of ``vacuum''. The energy given up by this term is the energy used to produce the particles. This example also allows us to study the possibility of symmetry restoration during preheating\\cite{lindekov,tkachev,kolbriotto}. While there have been arguments to the effect that the produced particles will contribute to the quantum fluctuations in such a way as to make the effective mass squared of the modes positive and thus restore the symmetry, we argued that they were unfounded. Whereas the effective squared mass oscillates, taking positive values during the early stages of the evolution, its asymptotic value is zero, compatible with Goldstone bosons as the asymptotic states. Furthermore, this says nothing about whether the symmetry is restored or not. This is signaled by the final value of the zero mode. In all the situations examined here, the zero mode is driven to a {\\em non-zero} final value. At this late time, the ``pions'' become massless, i.e. they truly are the Goldstone modes required by Goldstone's theorem. The arguments presented in favor of symmetry restoration rely heavily on the effective potential. We have made the point of showing explicitly why such a concept is completely irrelevant for the non-equilibrium dynamics when profuse particle production occurs and the evolution occurs in a highly excited, out of equilibrium state. Finally, we dealt with the issue of how to use our results to calculate the reheating temperature due to preheating in an inflationary universe scenario. Since our results are particular to Minkowski space, we need to assume that preheating and thermalization occur on time scales shorter than the expansion time, i.e. $H^{-1}$. We also need to assume that there is a separation of time scale between preheating and thermalization. Under these assumptions we can estimate the reheating temperature as $T_{reh} \\propto |M_R|$ in the case where the produced particles are massive and $T_{reh} \\propto |M_R|\\slash \\lambda^{\\frac{1}{4}}$ in the massless case. We have made the important observation that due to the large number of long-wavelength particles in the forbidden bands, a kinetic or Boltzmann equation approach to thermalization is {\\em inconsistent} here. A resummation akin to that of hard thermal loops, that consistently arises in the next order in $1/N$ must be employed. In equilibrium such a resummation shows that the scattering cross section for soft modes is perturbatively small despite their large occupation numbers. There is a great deal left to explore. Recently we reported on our study of the non-linear quantum field evolution in de Sitter and FRW backgrounds in refs.\\cite{desit} and \\cite{pfrw}, respectively. The next important step is to consider the background dynamics as a full backreaction problem, including the inflaton dynamics and the dynamics of the scale factor self-consistently. Such a detailed study will lead to a consistent and thorough understanding of the inflationary period, the post-inflationary and reheating periods. Further steps should certainly include trying to incorporate thermalization effects systematically within the $1\\slash N$ expansion. \\bigskip The preheating and reheating theory in inflationary cosmology is currently a very active area of research in fast development, with the potential for dramatically modifying the picture of the late stages of inflationary phase transitions. As remarked before, reliable and consistent estimates and field theory calculations have been done mostly assuming Minkowski spacetime. The matter state equations obtained in Minkowski\\cite{big} and recently in de Sitter backgrounds\\cite{desit} give an indication, through the Einstein-Friedmann equation, of the dynamics of scale factor and give a glimpse of the important physics to be unraveled by a deeper study. The formulation described in detail in these lectures are uniquely suited to provide complete description of the full dynamics of inflationary cosmology, from times prior to the phase transitions or the beginning of the chaotic era, through the inflationary regime, to the post-inflationary and reheating stage. Such a program provides the ultimate tool to test physical predictions of particle physics models. Thus this new consistent formulation provides the practical means to input a particle physics model and extract from it reliable {\\em dynamical} predictions which will have to ultimately be tested against the next generation of cosmological experiments. %" + }, + "9701/astro-ph9701094_arXiv.txt": { + "abstract": "We present phase-resolved spectroscopy of the bright, eclipsing polar \\hu\\ obtained with high time ($\\sim$30\\,sec) and spectral (1.6\\AA ) resolution when the system was in a high accretion state. The trailed spectrograms reveal clearly the presence of three different line components with different width and radial velocity variation. By means of Doppler tomography their origin could be located unequivocally (a) on the secondary star, (b) the ballistic part of the accretion stream (horizontal stream), and (c) the magnetically funnelled part of the stream (vertical stream). For the first time we were able to derive a (near-)complete map of the stream in a polar. We propose to use Doppler tomography of AM Herculis stars as a new tool for the mass determination of these binaries. This method, however, still needs to be calibrated by an independent method. The asymmetric light curve of the narrow emission line originating on the mass-donating companion star reveals evidence for significant shielding of 60\\% of the leading hemisphere by the gas between the two stars. ", + "introduction": "Polars or AM Herculis binaries are strongly magnetic cata\\-clysmic binaries. The white dwarf's magnetic field locks both stars in synchronous rotation and prevents the formation of an accretion disk. Instead mass transfer occurs via an accretion stream which is thought to follow initially a ballistic trajectory and later governed by the magnetic field, which leads the matter along magnetic field lines to the polar regions of the white dwarf. Most of the gravitational energy is released there, although some dissipative energy loss is expected to occur all along the stream, in particular in the coupling region where the stream is forced to leave the orbital plane. The stream manifests itself observationally in mainly three ways, (A) by so-called absorption dips observed in IR, optical or X-ray light curves (provided the inclination is high enough and the observer and the accretion spot are (blue in color) located on the same side of the orbital plane), (B) by residual emission during eclipses of the white dwarf, and (C) by strong, asymmetric and variable emission lines (H, He\\,{\\sc i}, He\\,{\\sc ii} and other high-ionized species) at optical and UV wavelengths (provided the system is in a high accretion state). The dips mentioned in (A) occur, if the line-of-sight crosses the out-of-plane parts of the stream. They are caused by photoelectric (at X-ray wavelengths), grey or free-free (at optical/IR wavelengths) absorption and are richly structured, indicating a highly fragmented stream. An overview of relevant data and analysis was given recently by Watson (1995). Phenomenon (B) is less studied because only a handful of eclipsing polars are known and the numbers of photons collected at eclipse which certainly arise from the stream are rather few, hence, model\\-ling of these events is widely missing in the literature, but nice observational data meanwhile have been obtained e.g.~for UZ For (Schmidt et al.~1993), WW Hor (Beuermann et al.~1990) and HU Aqr (Schwope et al.~1995). The emission lines (C) have attracted by far most attention because their occurence is one of the `classical' identification criterion. A diverse appearance is reported in the literature, they may consist of broad, narrow, high- and medium-velocity as well as (quasi-)stationary components, all with different radial velocity amplitudes and systemic velocities. These components may have a different appearance when a specific system is re-observed occasionally (see e.g.~the review of Mukai 1988). The mini\\-mum distinction usually made discerns between a narrow line attributed to the secondary star and a broad underlying component whose origin is assumed to be somewhere in the accretion stream (see e.g.~Liebert \\& Stockman 1985, who have shown that for the systems known at that time the radial velocity amplitude of the narrow line lies between that expected for the $L_1$ and the center of mass of the companion). Later on, Rosen et al.~(1987) have presented high-resolution spectra of V834\\,Cen and distinguished 4 subcomponents, which were thought to originate at different parts of the accretion stream and the secondary star. Schwope (1991) and Schwope et al.~(1991, 1993a) made use of repeated measurements of the narrow line in order to derive long-term ephemerides of the secondary star for the non-eclipsing systems QQ\\,Vul, MR\\,Ser and V834\\,Cen. Recent attempts to locate the line emission in polars applied tomographic techniques to data obtained on VV\\,Pup and RX\\,J0515.6+0105 (Diaz \\& Steiner 1994, Shafter et al.~1995). \\hu\\ is a system which has all of the observable features of the stream mentioned, at least occasionally. It was discovered independently by British and German astronomers during their optical identification programmes of ROSAT WFC and PSPC sources detected in the corresponding sky surveys (Schwope et al.~1993b, Hakala et al.~1993). Since it was the brightest eclipsing polar with the most extended eclipse at that time (meanwhile surpassed by RX\\,J0515.6+0105) it attracted immediate interest. We initiated an intensive optical follow-up study using 4 telescopes simultaneously and performed spectroscopy with high and low spectral resolution, high-speed photometry simultaneously in UBVRI-colours as well as standard one-channel photometry using a V-filter. All data have good signal/noise-characteristics and provide a comprehensive database suitable to reach deeper insight into the kinematics and accretion processes of the whole class. Here we present as a first paper in an upcoming row our results of high-resolution spectroscopy. ", + "conclusions": "We have found two distinct features in the trailed spectrograms of the four main emission lines in the blue spectral regime of \\hu, in particular the line of ionized Helium \\he2, which are potentially useful for a determination of the stellar masses. These are the amplitude of the radial velocity curve of the narrow emission line NEL and the location and curvature of the horizontal stream in the tomogram, the Doppler image of the HVC. An additional constraint on the mass ratio is set by the observed eclipse length. All these constraints are shown in a $(Q, i)-$plane in Fig.~\\ref{qi_rel} with $Q$ being the mass ratio and $i$ the orbital inclination. The eclipse length is dependent on the size of the secondary star's Roche lobe and the inclination, hence relation (1) in the diagram is purely geometrical (it is of course implicitly assumed that the secondary fills exactly its Roche lobe) (Chanan et al.~1976). The solid line is valid for the measured eclipse length, the nearby dashed line is valid for the maximum correction which accounts for the offset of the accretion spot from the white dwarfs center, see above. Relation (2) in the diagram for the NEL makes use of a factor which relates the center of light of the illuminated hemisphere to the center of mass. This factor is dependent on $Q$ and shown in Horne \\& Schneider (1989) for the optically thin and in Schwope et al.~(1993) for the optically thick models. The differences between both modes are negligible as far as this factor is concerned. The computation of the curves (2) in the $(Q,i)-$plane makes use further of the Roche geometry (approximation of the spherical equivalent Roche-radius as a function of $Q$ given by Eggleton (1983)) and Kepler's third law. We have assumed that the secondary is a main sequence star and can be described by a ZAMS mass-radius relation for late-type stars. With $M_2 = 0.146$\\,M$_\\odot$, the empirical relation by Caillault \\& Patterson (1990) predicts the least compact secondary for \\hu, the other extreme with $M_2 = 0.206$\\,M$_\\odot$ is due to the relation by VandenBerg et al.~(1983) (solid and dashed lines in Fig.~\\ref{qi_rel}). Relation (3) finally is the result of a fit-per-eye to the location and curvature of the horizontal stream in the tomogram which relies on the assumption that the center of light along the ridge of \\he2 \\ marks the single-particle trajectory. This results in the smallest estimates of $Q$ with $Q \\simeq 2.5$. \\begin{figure}[t] \\psfig{figure=qi_rel_ps,width=88mm,bbllx=8mm,bburx=200mm,bblly=53mm,bbury=245mm,clip=} \\caption[qi_rel]{\\label{qi_rel} Mass determination of \\hu \\ using the observed length of the eclipse (1), the radial velocity amplitude of the illuminated hemisphere of the secondary star with different mass-radius relations (2), and the curvature of the ballistic stream in the Doppler map (3). } \\end{figure} It is clear from Fig.~\\ref{qi_rel} (as well as from several previous Figures) that no unique solution exists presently in the framework of the models presented here. Relations (1) and (2) suggest high values of the mass ratio and the inclination with $Q \\simeq 4.5 - 5.5, i > 85\\degr$ for the Helium lines and $Q = 5 - 5.5, i > 87\\degr$ for the Balmer lines (the latter are not shown in Fig.~\\ref{qi_rel}). The curvature of the stream of \\he2, \\h1\\ and H$\\beta$ on the other hand suggests, that the mass ratio might be as low as $Q = 2.5$ and that the inclination correspondingly might be as low as $i \\simeq 80\\degr$. Both estimates of the mass ratio seem to exclude each other presently and it appears likely that neither the NEL (in combination with $K_2$-correction and ZAMS-approximation) nor the HVC (fitted by a single-particle trajectory) can be used straightly to give an accurate mass estimate. The disagreement between both approaches becomes smaller if one assumes that NEL emission is more concentrated towards the center of mass of the secondary than we have assumed. An absorbing cloud of matter in the $L_1$-region which would suppress emission from the `nose' of the secondary could e.g.~provide the necessary bias of the effective radial velocity amplitude towards higher values (and could also lower the width of the line profiles, Fig.~\\ref{fwhm_nel}). The expected re-emission might be responsible for the longish extension of the bright spot near the $L_1$-point in the tomogram of H$\\beta$. That such bias exists appears likely from the observed higher NEL-velocity of the Hydrogen Balmer lines with respect to the Helium lines. The `nose' is more effectively shielded in the more optically thick Hydrogen lines. The best way to solve the conflict and to calibrate the M/R-relation of the secondary star would be to measure the radial velocity of the secondary star independently using e.g.~photospheric absorption lines. This will help to sort out which feature (NEL, HVC) of which line may serve as {\\it the} estimator of the mass ratio. It requires \\hu\\ to be observed in a low state of accretion. Such a measurement is highly demanded in order to calibrate either of the methods used here for mass determination. Another hint to the likely mass of the stars and the mass ratio could be derived from ingress/egress length into/from eclipse of the white dwarf itself, rather than ingress/egress of the accretion spot on it. Our model with $Q = 2.5$ implies that the secondary star in \\hu\\ must be as massive as 0.35\\,M$_\\odot$. It would be much more compact than a ZAMS-star and in consequence much more luminous. This would imply that the distance estimates given by Glenn et al.~(1994) and Schwope et al.~(1993) are lower limits only. \\medskip We have presented high-resolution spectral observations of the eclipsing AM Herculis star \\hu\\ obtained when the system was in a high state of accretion. We have identified three emission line components in our trailed spectra of \\he2, \\h1, and the H-Balmer lines. These could be uniquely identified as originating on the illuminated hemisphere of the secondary star (the narrow emission line NEL), the stream in the orbital plane (the high-velocity component HVC) following a ballistic trajectory, and the stream coupled onto field lines out of the orbital plane (the broad-base component BBC). We have uniquely identified the horizontal stream in this AM Her-binary by means of Doppler tomography and could trace it from the $L_1$-point down to the stagnation region. On the assumption that the NEL of \\he2 is formed on the surface of the Roche-lobe of the companion star we have determined its mass. This estimate is in disagreement with an estimate derived from the tomogram. We propose that Doppler tomography of the ballistic stream can be used as a new tool for the mass determination of accreting binaries provided we can calibrate this method using an independent mass determination." + }, + "9701/astro-ph9701211_arXiv.txt": { + "abstract": "Galaxy morphologies in clusters have undergone a remarkable transition over the past several billion years. Distant clusters at $z \\sim 0.4$ are filled with small spiral galaxies, many of which are disturbed and show evidence of multiple bursts of star-formation. This population is absent from nearby clusters where spheroidals comprise the faint end of the luminosity function. Our numerical simulations follow the evolution of disk galaxies in a rich cluster owing to encounters with brighter galaxies and the cluster's tidal field---galaxy harassment. After a bursting transient phase, they undergo a complete morphological transformation from ``disks\" to``spheroidals\". We examine the remnants and find support for our theory in detailed comparisons of the photometry and kinematics of the spheroidal galaxies in clusters. Our model naturally accounts for the intermediate age stellar population seen in these spheroidals as well as the trend in dwarf to giant ratio with cluster richness. The final shapes are typically prolate and are flattened primarily by velocity anisotropy. Their mass to light ratios are in the range 3---8 in good agreement with observations. ", + "introduction": "The origin of the Hubble sequence (Hubble 1936, Sandage 1961) remains a long standing puzzle in astronomy. Giant galaxies range from ellipticals, slowly-rotating dense spheroids with little gas, to late--type spirals that are rapidly rotating thin disks of gas and stars. In recent years, two distinct classes of ellipticals have been recognized that are easily separated in plots of nearly any two of their properties, such as central surface brightness versus luminosity (Ferguson and Binggeli 1994, Kormendy 1985). One class includes the bright, giant ellipticals and extends to the rare ``dwarf ellipticals\" with high surface brightness, M32 being the prototype. The second class consists of low surface brightness spheroidal galaxies that have luminosities $M_{_B} \\gsim -17$, about 3 magnitudes fainter than \\l*, the characteristic break in the luminosity function. (Throughout this paper, all distance dependent quantities assume a Hubble constant of $100h \\kms Mpc^{-1}$, with $h=1$.) The dwarf spheroidal galaxies (dSph) in our Local Group of galaxies with magnitudes in the range $ -8 \\gsim M_B \\gsim -12$ are often considered to be the low luminosity extreme of this sequence, but nearly all other known galaxies in this class reside in clusters (Vader and Sandage 1991). This second class is referred to as ``dwarf ellipticals'' by Ferguson and Binggeli (1994) and ``spheroidals'' by Kormendy (1985) who reserves the designation ``dwarf ellipticals'' for galaxies like M32. After attempting a clumsy hybrid of labels to stay clear of this confusion, we reluctantly follow Kormendy. There is no shortage of theories for the formation of giant galaxies and the origin of their Hubble sequence. Hoyle (1945) calculated an upper limit to the mass that can radiatively cool and proposed that this imprints a characteristic galaxian mass scale. Gott (1977) and Gunn (1982) promoted two epochs of galaxy formation with ellipticals and bulges forming in an early epoch of Compton cooling and spirals forming by radiative cooling at a later time. This model has faded from popularity because the two formation epochs means that morphological types must arise from fluctuations with vastly different amplitudes, must cool by different physical processes yet have comparable velocity scales and masses. Most theories for the origin of the Hubble sequence follow Jeans' (1938) assertion that angular momentum is the control parameter. Among the giant galaxies, the Hubble sequence can also be viewed as a mass sequence (Tully, Mould and Aaronson 1982) or virial velocity sequence (Meisels 1983). The larger ellipticals naturally occur in clusters owing to statistical biasing in the hierarchical gravitational collapse model (Bardeen \\etal 1986), where larger peaks are found to be more strongly clustered than smaller ones. Unfortunately, the variation in the angular momentum of protogalaxies in the hierarchical model is small compared to the observed variation across the Hubble sequence. Furthermore, the angular momentum of halos correlates poorly with any other property such as local overdensity (Barnes and Efstathiou 1987, Ryden 1988). In order to match the wide variations in the observed properties of galaxies we need a mechanism that can increase the available stretch and correlations in the hierarchical model. One way to do this is to form elliptical galaxies by the merger of two spirals (Toomre 1977). Despite continuous assertions of its demise (Ostriker 1980, van den Bergh 1982, Carlberg 1986), it seems fairly robust (Lake and Dressler 1986, Lake 1989, Zepf and Ashman 1993). However, ellipticals are generally found within larger structures with collapse times that are shorter than the cooling time required to make a disk. Lake \\& Carlberg (1988ab) and Katz (1991) found that the Hubble sequence could also arise if the relative amplitude of small scale fluctuations increased with mass or binding energy. If larger mass fluctuations have greater small scale power, then their lumps lead to a transfer of angular momentum from the luminous material to the dark matter. An elliptical galaxy easily results from a lumpy collapse model without disks having already formed. Similarly, Lake \\& Carlberg noted that small scale power must be suppressed in order for a disk to survive, an effect that was reexamined by Toth and Ostriker (1992). Spheroidal formation theories have a shorter history and have focused on incremental changes to the giant galaxy theory. The popular model of Dekel and Silk (1986) combines the notion of lower amplitude peaks in the hierarchical model with the use of stellar winds or supernovae to reshape the small galaxies that have low binding energy by expelling their gas (Larson 1974, Vader 1986). There are many problems with this scenario, {\\it e.g.} the clustering properties of dwarfs is opposite to the expectations of the Dekel and Silk model (Ferguson and Binggeli 1994). The model also has the seemingly impossible chore of explaining the general properties of both rapidly-rotating gas-rich dwarfs and gas free dwarf spheroidals. Recent Hubble Space Telescope (HST) observations (Dressler \\etal 1994a) have revealed that the morphologies of galaxies in clusters have changed dramatically since $z \\sim 0.4$. Over 20 years ago, Butcher and Oemler (1978, 1984) discovered a large population of ``blue galaxies\" in clusters at $z\\sim 0.4$. The HST frames show that at $z \\sim 0.4$, the giant ellipticals are already in place, but the ubiquitous ``blue galaxies\" are distorted spirals that have vanished from clusters at the present epoch (Dressler \\etal 1994a). The population difference is greatest at lower luminosities; 90\\% of galaxies fainter than $L_*/5$ in distant clusters are bulgeless ``Sd'' disk systems, whereas 90\\% are spheroidals in nearby clusters such as Virgo or Coma (Sandage \\etal 1985, Thompson and Gregory 1993). Couch \\etal (1994) present spectroscopic evidence that the distorted blue galaxies at $z \\sim 0.3$ have undergone multiple burst events separated by 1--2 Gyr. In hierarchical clustering models, the influx of field galaxies into clusters peaks at $z \\sim 0.4$ (Kauffmann 1995). A complete picture emerges of both the Butcher--Oemler effect and the morphological transformation if we can identify a mechanism that begins when spiral galaxies enter clusters, causes a few episodes of distortion and star formation within each galaxy over an interval of 1--2 Gyr and ultimately turns the dwarf spirals into spheroidals in 4--5 Gyr. Several possible mechanisms have been proposed for the Butcher--Oemler effect: mergers (Icke 1985, Miller 1988), gas compression in the cluster environment (Dressler and Gunn 1983) and tidal compression by the cluster (Byrd and Valtonen 1990, Valluri 1993). Nearly all of these mechanisms can make starbursts, but they all fail other critical tests (c.f. Valluri and Jog 1993). Merging is a one time event in contrast to the spectroscopic evidence of multiple bursts found by Couch \\etal (1994). From their HST images, Oemler \\etal (1996) conclude that merging is implausible as the blue galaxy fraction is large and the merging probability is low. They also observe disturbed spirals throughout the cluster, whereas both ram pressure stripping and global tides only operate efficiently near the cluster's center. Furthermore, Valluri and Jog (1993) have shown that the observed relationship between HI deficiency and size is exactly opposite to that expected from ram pressure or evaporative models. Finally, there is no correlation between the spiral fraction and X-ray luminosity of clusters. None of the other models for the Butcher--Oemler effect discussed above address the issue of present day remnants: how are the ubiquitous small spheroidals created as the distorted spirals disappear? Recently, we proposed a new mechanism for the Butcher--Oemler effect in clusters---``galaxy harassment\" (Moore \\etal 1996). At speeds of several thousand kilometers per second, close encounters with bright galaxies cause impulsive gravitational shocks that can severely damage the fragile disks of Sc--Sd galaxies. We will show that these collisions are frequent enough that harrasment occurs throughout the cluster. Moore \\etal (1996) used numerical simulations to compare harassed galaxies to HST frames of galaxies in clusters at $z \\gsim 0.3$. They stated that the cumulative effect of such encounters changes a disk galaxy into a spheroidal galaxy, thus identifying the present-day remnants of the disturbed blue galaxies and explaining the change in galaxy morphologies in clusters since $z \\sim 0.4$. In this paper, we provide detailed comparisons of the harassed remnants of late type spirals with the photometric and kinematical properties of dwarf spheroidal galaxies. ", + "conclusions": "We have followed the evolution of small disk galaxies in clusters using realistic simulations of the combined action of rapid encounters with cluster galaxies and global tidal heating---``galaxy harassment\". We find that the first shocks create the population of distorted spiral galaxies seen in HST images $z \\sim 0.4$ cluster. Over a period of a few Gyr, these galaxies evolve into remnants that have surface density profiles, shapes, kinematics, stellar populations and clustering properties that qualitatively match the spheroidal systems in present-day clusters. The assumptions and parameters of our theory are minimal. It relies only on gravitational interactions with the global cluster potential and other galaxies. The strength of these effects is well determined by the observed properties of the clusters and their brighter members. Non gravitational effects such as ram pressure stripping, wind driven expulsion of gas or other hydrodynamic interactions can speed the timescale for the evolution of the gas, but will not sufficiently alter the stellar disk of the galaxies observed at $z \\sim 0.4$. Further, the harassment timescale matches the interval of multiple starbursts that is inferred in clusters at $z \\sim 0.4$ (Couch \\etal 1994). This provides a limit on the magnitude of additional hydrodynamic effects, assuring us that harassment must be the dominant mechanism. In this paper, we have focused on the remnants. However, it is clear that roughly 50\\% of the stars escape into the intracluster medium. They do so within debris tidal streams and thin tails that gradually disperse. A study of these low surface brightness features and the general diffuse light in clusters will be the subject of a future paper." + }, + "9701/astro-ph9701027_arXiv.txt": { + "abstract": "{ We report continuous monitoring of Cygnus X-1 in the 1.3 -- 200 keV band using ASM/RXTE and BATSE/CGRO for about 200 days from 1996 February 21 to 1996 early September. During this period Cygnus X-1 experienced a hard-to-soft and then a soft-to-hard state transition. The low-energy X-ray (1.3-12 keV) and high-energy X-ray (20-200 keV) fluxes are strongly anti-correlated during this period. During the state transitions flux variations of about a factor of 5 and 15 were seen in the 1.3-3.0 keV and 100-200 keV bands, respectively, while the average 4.8-12 keV flux remains almost unchanged. The net effect of this pivoting is that the total 1.3-200 keV luminosity remained unchanged to within $\\sim$15\\%. The bolometric luminosity in the soft state may be as high as 50-70\\% above the hard state luminosity, after color corrections for the luminosity below 1.3 keV. The blackbody component flux and temperature increase in the soft state is probably caused by a combination of the optically thick disk mass accretion rate increase and a decrease of the inner disk radius.} ", + "introduction": "Cygnus~X-1 is one of the brightest high-energy sources in the sky, with an average 1-200 keV energy flux of $\\sim$3$\\times$10$^{-8}$ erg/cm$^{2}$/sec. Most of the time its X-ray spectrum is very hard, with a high-energy ($> 20$ keV) X-ray flux around or above that of the Crab Nebula while its low-energy ($<10$ keV) X-ray flux is around 0.5 Crab. Occasionally, the spectrum of Cyg~X-1 becomes much softer; then its low-energy X-ray flux increases to 1 -- 2 Crab while the hard X-ray flux decreases to 0.5 Crab or less. On the basis of the radial-velocity curve of its O9.7 Iab companion star, Webster and Murdin (1971) and Bolton (1971) concluded that the compact star in Cyg~X-1 may be a black hole (BH); their results have been confirmed by later detailed work of Gies and Bolton (1986) who concluded that the mass of the compact object in Cyg~X-1 is greater than 7 \\msol, but more probably 16 \\msol, far exceeding the theoretical (and observational) upper mass limit of 3.2 \\msol\\ for a neutron star. It is thus the first stellar mass BH candidate (BHC) (cf. Liang and Nolan 1984 and Lewin and Tanaka 1995, for reviews and references therein). Cyg~X-1 is often considered as the canonical BHC; many of its characteristics, such as its high X-ray luminosity above 100 keV, the ultra-soft component in its X-ray spectrum, the hard/soft X-ray flux anti-correlation, and the rapid X-ray flux variability, are shared by other systems, which on the basis of a dynamical mass estimate may contain a BH. It is, however, currently not understood why accreting BHs show these characteristics. In spite of extensive studies over the last three decades, the mass accretion conditions near the central compact object in Cyg~X-1 and other BH systems, and the high-energy radiation mechanism are still poorly understood. This remains an outstanding issue in high-energy astrophysics. State transitions fully covered over a large range in X-ray photon energy may provide us with important clues towards a better understanding of these BH X-ray binaries. During a transition, a rich collection of information may be obtained, such as flux variations at many time scales in all energy bands, energy spectral evolution, correlations between different energy bands, etc; such data should provide useful tests of various theories and models. Previously, several transitions between the hard state (HS) and soft state (SS) of Cyg~X-1 have been observed. In 1971, a soft-to hard (S-to-H) state transition was observed in the 2-20 keV band (Tananbaum \\etal\\ 1972). In 1975, a complete transition was observed with {\\it Ariel} V between 3 and 6 keV (Holt \\etal\\ 1976), and by {\\it Vela} between 3 and 12 keV (Priedhorsky \\etal\\ 1983). These observations were limited to rather low X-ray energies. The only broad-band observations of a transition were obtained in 1980, with {\\it Hakucho} between 1 and 12 keV (Ogawara \\etal\\ 1981) and with HEAO-3 between 48 and 183 keV (Ling \\etal\\ 1983). However, the SS onset was observed only with HEAO-3 as a rapid decrease in the 48-183 keV flux, and then later the SS was observed with {\\it Hakucho} only when the 1-12 keV flux was already a factor of 2-3 higher than its HS level before the transition. Therefore, no simultaneous low- and high-energy X-ray observations during S-to-H or hard-to-soft (H-to-S) transitions of Cyg~X-1 have been made so far. In this paper we report results obtained from the near continuous monitoring of Cyg~X-1 during a complete state transition from the HS to the SS (Cui 1996; Cui, Focke and Swank 1996; Zhang \\etal\\ 1996a), and then back to the HS (Zhang \\etal\\ 1996b, 1996c), observed simultaneously with ASM/RXTE (1.3-12 keV) and BATSE/CGRO (20-200 keV) for about 200 days from February to September 1996. Our results are complemented by detailed pointing observations with ASCA (0.5-10 keV) (Dotani \\etal\\ 1996), and with the PCA and HEXTE (2-250 keV) aboard RXTE (Cui, Focke zhang Swank 1996; Belloni \\etal\\ 1996; Cui \\etal\\ 1997), and with SAX (0.1-300 keV) (Piro \\etal\\ 1996) and with OSSE/CGRO (50-600 keV) (Phlips \\etal\\ 1997). We will focus on variations in the total luminosity during this period and on the correlations between the low-energy and high-energy X-ray fluxes. ", + "conclusions": "Compared to the high-energy observations of the previous Cyg~X-1 SS transitions, the data presented here are the first set covering the whole transition episode continuously over a broad energy band. The observed 1.3-15 keV ({\\it Ariel} V) spectral evolution (spectral pivoting and power law index changes) (Chiappetti \\etal\\ 1981) during the 1975 S-to-H state transition is quite similar to that presented here. The initial HXR flux and spectral steepening observed with HEAO 3 between 48-183 keV (Ling \\etal\\ 1983) and the soft X-ray flux level between 1-12 keV observed later with {\\it Hakucho} (Ogawara \\etal\\ 1981) are similar to that during the SS presented here. Both previous SS lasted between 60-100 days, again similar to the 1996 SS duration. So this 1996 SS transition is qualitatively similar to the previous ones. It is thus reasonable to assume that the same physical mechanism is responsible for all of them. GX~339-4 is the only other BHC observed to have recurrent state transitions. From a comparison with its 1981 H-to-S state transition (Maejima \\etal\\ 1984), we find that the HS spectra of this source and Cyg~X-1 are quite similar. The ratio between the SS and HS total luminosity observed from GX~339-4 (e.g., Ricketts 1983) is also similar to that of Cyg~X-1 presented here. During a SS observation of GX~339-4 in 1983 (Makishima \\etal\\ 1986), the HXR flux increased significantly while the soft X-ray flux remained nearly unchanged, also similar to the HXR flux increase in the second half of the SS of Cyg~X-1. The photon index of the power law changed from --0.9 to --2.1 following the HXR flux increase of GX~339-4, different from the near constant value of --2.5 during much of the SS in Cyg~X-1. The power law tail during a `very high' state of GX~339-4 (Miyamoto \\etal\\ 1991), when the overall flux was about a factor of 2-3 higher than during the previous SS observations, is however, very similar to that of Cyg~X-1 in the SS. Qualitatively similar S-to-H state transitions have also been observed from the low-mass X-ray binary BH system GS~1124--683 (Ebisawa \\etal\\ 1994) and the neutron star system 4U~1608--52 (Mitsuda \\etal\\ 1989). Therefore similar physical mechanisms might operate in all of them. The `hard' and `soft' states we refer to in this paper are usually called the `low' state, and the `high' (and `very high') states. This is due to the fact that the early observations of them were made in the low-energy X-ray band ($<$20 keV; see e.g., Tananbaum et al. 1972), so the terms `low' and `high' refer to the low and high values of the low-energy X-ray fluxes, respectively. It is generally believed that the low-energy X-ray flux tracks the total mass accretion rate of the system; therefore, these `low' and `high' states have been considered to correspond to low and high values of \\mdot~, respectively. This idea gained strong support in the unified scheme of source states of both neutron star and BH X-ray binaries proposed by Van der Klis (1995). This picture may be incomplete in view of the lack of a strong variation of the total luminosity during the state transitions which we observed in 1996. The blackbody component detected in both the hard and soft states is generally thought to originate from an optically thick and geometrically thin accretion disk near the BH (see, e.g., Mitsuda et al. 1984). According to current models the hard power law spectral component is likely produced in a very hot and optically thin region, through Comptonization of low-energy X-ray photons; the detailed nature of the optically thin region and the mechanism that makes it very hot distinguish these different models. The near constant luminosity during the transitions indicates that they are driven by a redistribution of the gravitational energy release between the optically thick and the optically thin regions. Therefore {\\it a H-to-S state transition, and its reverse, probably reflects a change in the relative importance of the energy release in the optically thin and thick regions of the accretion disk near the BH and this may not require a substantial change in the total accretion rate.} Spectra emitted by accretion disks around black holes are well described by the `multi-temperature disk blackbody model' of Mitsuda et al. (1984), which has as fit parameters the inner disk radius, $R_{\\rm in}$, and the temperature, $T_{\\rm in}$, at that radius. It is usually assumed that the inner disk radius equals three times the Schwarzschild radius, and this provides acceptable mass estimates for the black hole (see Tanaka and Lewin 1995, for a review). The blackbody fits made to the high-energy tail of this multi-temperature disk BB model model provide good fits for a blackbody temperature, $T_{\\rm bb}$, which turns out to be equal to the temperature in the disk at a radius of about 7 Schwarzschild radii (Ross et al. 1992); correspondingly, one has $T_{\\rm bb} \\simeq 0.7 T_{\\rm in}$. Independent of the ratio of inner disk radius to the Schwarzschild radius, one has $T_{\\rm bb} \\propto T_{\\rm in}$. Since the multi-temperature disk model the bolometric luminosity, ($L_{\\rm bol,disk}$) follows the proportionality relation $L_{\\rm bol,disk} \\propto R_{\\rm in}^2~T_{\\rm in}^4$, one has $(L_{\\rm h}/L_{\\rm s}) = (R_{\\rm h}/R_{\\rm s})^2 (T_{\\rm h}/T_{\\rm s})^4$, by applying the proportionality relation in both the HS and SS, where the subscripts `h' and `s' denote the hard and soft states respectively. From the observations we infer that $L_{\\rm s}/L_{\\rm h} \\la 6$, and $T_{\\rm s}/T_{\\rm h} \\ga $ 2.1 to 2.8. Consequently, we obtain $R_{\\rm h}/R_{\\rm s} \\ga $ 1.8 to 3.2. This result suggests that {\\it during the H-to-S transition the inner radius of the geometrically thin and optically thick disk changed from $\\ga 170$ km to $\\sim 70$ km} (the latter value has been obtained by Dotani et al. 1997, from ASCA observations during the soft state). Applying the relationship between the mass accretion rate, inner disk radius and the inner disk temperature in the multi-temperature disk model we have $\\dot M_{\\rm s}/\\dot M_{\\rm h} \\la $ 2.0 to 3.4. This change in inner disk radius accompanying the spectral state changes suggests a picture in which during the hard state advection of internal disk energy into the black hole, as proposed by Narayan (1996), dominates within a radial distance of $\\ga 170$ km from the hole. In the soft state the flow in the inner disk may still be advection dominated, with only a moderate fraction of the mass flow passing through an optically thick inner disk. Variations of the inner disk radius have also been discussed by Ebisawa \\etal\\ (1996), in the context of the mass accretion and high-energy radiation model of Chakrabarti and Titarchuk (1995), and have recently been proposed by Belloni et al. (1996) as an explanation for the rapid variability of the black-hole candidate GRS~1915+105 (albeit at much higher mass accretion rates). A potential problem for any model that purports to explain the spectral transitions in terms of an instability in the inner disk region is that the transitions seem to be accompanied by the gradual changes in the slope and the flux of the HXR power law component, which occur over a very long time scale ($\\sim$ weeks). In fact, the 1996 S-to-H transition was predicted from the HXR flux increase (Zhang et al. 1996c). It is possible that Cyg~X-1 underwent a very slow change in the mass accretion rate, which initially only affected the properties of the hard X-ray emission region. The sudden soft (1.3-3.0 keV) X-ray flux increase and decrease over a much shorter time scale, during the state transitions, may then represent the crossing of a threshold, at which the radiative efficiency exceeds a value required for the formation of an optically thick disk down to the innermost stable orbit. In view of the observational limitations the above remarks are necessarily somewhat speculative. It would appear that a better understanding of the cause of the spectral state changes in Cyg X-1 requires that further monitoring of the long-term behaviour of Cyg X-1 include low-energy ($\\leq 1$ keV) coverage of the source as well." + }, + "9701/astro-ph9701071_arXiv.txt": { + "abstract": "Sequences of infinitesimally thin, uniformly rotating, self-gravitating relativistic discs with internal two-dimensional pressure have been constructed. It is shown that in weaker relativistic configurations the sequences undergo a continuous bifurcation from a disc to a ring structure, while in stronger relativistic cases the sequences terminate at the mass-shed limit where gravitational forces are exactly balanced by centrifugal forces. ", + "introduction": "Relativistic, self-gravitating discs may have played an important role during the process of galaxy formation in the early universe. For example, they can be considered as possible predecessors of galaxies today that contain a massive black hole in their centre. Also, during the final evolution of a close binary system, consisting of two neutron stars, the orbit shrinks rapidly due to the emission of gravitational waves and the two stars merge finally into one object. Shortly after this coalescence, the system can be approximately described as a highly flattened disc-like structure which may display relativistic effects. These systems may then be modeled, in a simplified approach, by infinitesimally thin discs where the vertical variation has been averaged onto the midplane. Uniformly rotating relativistic discs with zero internal pressure have been studied in detail by Bardeen \\& Wagoner (1971). Here an expansion method was used in which the problem was reduced to a set of ordinary differential equations, which were then solved numerically for the higher order corrections. Nevertheless, using this method, strongly relativistic disc structures could also be calculated accurately. Recently, the analytical solution of this problem in terms of integral equations was obtained (Neugebauer \\& Meinel 1993), which could subsequently be solved utilizing hyper-elliptic functions (Neugebauer \\& Meinel 1995). This solution represents the first analytic solution of a continuously extended rotating object in general relativity. In these approaches a given equation of state ($p=0$) and rotation law ($\\Omega=$const.) are assumed, and the Einstein-equations are solved directly. The method of mirror images is another frequently used option to construct models of relativistic discs. Here one begins with a known axisymmetric metric (e.g. Kerr), makes a cut in the upper half of the $\\rho-z$ plane where $\\rho$ and $z$ are cylindrical coordinates, and identifies the metric functions there with those of a symmetric cut in the lower half which joins the first at the $z=0$ plane (see Pichon \\& Lynden-Bell 1996, and references therein). This method yields disc solutions where the corresponding matter properties, i.e. the energy-momentum surface distribution, must be inferred from the pasted solution after the two cuts have been identified. The method was first used for relativistic discs by Turakulov (1990). It is known that, in the Newtonian limit, the uniformly rotating pressure-less dust disc is unstable to all perturbing modes (Binney \\& Tremaine 1987), a result which most likely carries over into the relativistic regime (Bardeen \\& Wagoner 1971). The addition of pressure tends to stabilize the disc. For Newtonian discs it was shown recently (Kley 1996) that uniformly rotating, pressurized discs undergo a continuous bifurcation process into a ring-like structure via a sequence of dumb-bell shaped equilibria. This process has its three-dimensional analogue in the bifurcation of rotating stars into rings (Eriguchi \\& Sugimoto 1981). In this paper we generalize the relativistic zero-pressure disc solutions by imposing a given (polytropic) equation of state relating the surface pressure with the surface density. This pressure may be thought of as being generated by the random internal peculiar velocities in a galaxy, and it has a stabilizing effect. This work on the other hand also generalizes the Newtonian results on thin Maclaurin discs and their bifurcations to the relativistic case. In Section 2, the relevant equations are given and the numerical method is outlined in Section 3. In Section 4 the results of the numerical computations are presented and in Section 5 we conclude. ", + "conclusions": "We calculated the structure of uniformly rotating, infinitesimally thin, relativistic self-gravitating discs with internal pressure. The pressure is given by a polytropic equation of state. The polytropic exponent $\\gamma=3$ was used, since in this case there exists an exact solution in the Newtonian limit. This special choice corresponds to three-dimensional bodies of constant densities as well (see Hunter 1972). Thus, we may make a comparison of the results found here for flat discs with those of rotating, homogeneous relativistic stars. As Butterworth \\& Ipser (1976) showed, sequences of homogeneous, rotating relativistic stars usually terminate at the mass-shed limit. Close to the Newtonian case however, they were not able to follow the sequences to this limit. We suggest that, as weaker relativistic discs bifurcate into a ring, weaker relativistic stars of constant density bifurcate into a toroid structure as well, and do not end at the mass-shed limit. In the Newtonian limit, the structure of these constant density toroids were calculated by Eriguchi \\& Sugimoto (1981). The pressure-less discs possess ergo-regions, where the dragging of inertial frames would force observers to rotate. These first appear at a single point within the disc at $z\\subscr{e}=1.41$ and reach the edge of the disc at $z=1.89$ (Meinel \\& Kleinw\\\"achter 1993). They also occur in rotating stars (Butterworth \\& Ipser 1976). For discs with internal pressure however, we found no indication of the existence of ergo-regions. It is to be expected that there exists a continuous transition from the zero-pressure $K=0$ case to discs with non-vanishing pressure ($K$ small). However, this connection could not be demonstrated numerically, as disc sequences end either at the mass-shed limit or bifurcate into rings before reaching the pressure-less limit. One might speculate, perhaps, that such a connection can be achieved by a disc which consists of two parts: a pressure supported central region ($0 < \\rho < \\rho\\subscr{p}$), surrounded by a dust disc ($\\rho\\subscr{p} < \\rho < \\rho\\subscr{d}$). The dust disc would refer to $\\rho\\subscr{p}=0$, and the disc sequences to $\\rho\\subscr{p}=\\rho\\subscr{d}$. In some rotating relativistic stars, there are sequences of supra-massive stars (Cook, Shapiro \\& Teukolsky 1992), which are so massive that they exceed the rest mass of a non-rotating star, and can only exist for non-zero rotation. In the case of flat relativistic discs with non-zero internal pressure, we did not find any supra-massive disc sequences. Possibly, these negative results are a result of the particular equation of state where in the strong relativistic limit the mass is always concentrated in the centre of the disc, as shown in Fig.~4 for the non-rotating case. Other relations between pressure and surface density, in particular a smaller $\\gamma$, could also possibly lead to supra-massive sequences or ergo-regions. The numerical method developed here is sufficiently general to be applied to the study of differentially rotating discs, and can easily be extended to three-dimensional rotating stars as well. This may be the subject of a future further investigation. A stability analysis of the computed equilibrium configurations lies beyond the scope of the present paper." + }, + "9701/astro-ph9701159_arXiv.txt": { + "abstract": "The spiral structure of the low surface brightness galaxies F568-6 (Malin 2) and UGC~6614 is large scale, with arms that wrap more than half a revolution, and extend out to 50 and 80 kpc in UGC~6614 and F568-6 respectively. The density contrasts observed in the \\ion{H}{1} maps are high, with arm/interarm contrasts of $\\sim$2:1, whereas the velocity perturbations due to spiral structure are low, in the range 10--20 km/s and 10--30 km/s in UGC~6614 and F568-6 respectively. Upper limits for the disk mass-to-light ratios are estimated by considering the minimum velocity perturbations in the \\ion{H}{1} velocity field that should result from the spiral structure observed in the $R$ band images. The weak observed response in the $\\phi$ velocity component limits the mass-to-light ratios of the disk inside a scale length to $M/L \\lta 3$ and 6 for UGC~6614 for F568-6 respectively (in solar units) based upon azimuthal variations observed in the $R$ band images. These limits are sufficiently strong to require a significant dark matter component even in the central regions of these galaxies. Our limits furthermore imply that this dark matter component cannot be in the form of a cold disk since a cold disk would necessarily be involved in the spiral structure. However, a more massive disk could be consistent with the observations because of a non-linear gas response or if the gas is driven by bar-like distortions instead of spiral structure. To produce the large observed arm/interarm \\ion{H}{1} density variations it is likely that the spiral arm potential perturbation is sufficiently strong to produce shocks in the gas. For a forcing that is greater than $2\\%$ of the axisymmetric force, $M/L \\gta 1$ is required in both galaxies in the outer regions. This is equivalent to a disk surface density between $r =$ 60--120\\arcsec\\ in UGC~6614 of 2.6--1.0 $M_\\odot/{\\rm pc}^2$ and between $r =$ 40--90\\arcsec\\ in F568-6 of 6.6--1.0 $M_\\odot/{\\rm pc}^2$ assuming that the amplitude of the variations in the disk mass is the same as that observed in the $R$ band. These lower limits imply that the stellar surface density is at least of the same order as the gas surface density. This is consistent with the large scale morphology of the spiral structure, and the stability of the gas disk, both which suggest that a moderate stellar component is required to produce the observed spiral structure. ", + "introduction": "UGC~6614 and F568-6 were observed in $R$ band and in \\ion{H}{1} by \\cite{pic97} to investigate their neutral hydrogen and kinematic properties. These data are displayed as overlays in Figures 1 and 2 and are described in detail in \\cite{pic97}. Table 1 lists some observation parameters and basic properties from \\cite{pic97}. UGC~6614 and F568-6 are both low surface brightness galaxies with central surface brightnesses several magnitudes fainter than sky level ($\\mu_{R}(0) = 22.9$ and 22.1 for UGC~6614 and F568-6 respectively). Their disk exponential scale lengths are large (14 and 18 kpc for UGC~6614 and F568-6 respectively) and they both contain copious amounts of \\ion{H}{1} ($2.5$ and $3.6 \\times 10^{10} M_\\odot$ for the two galaxies respectively, \\cite{pic97}). We assume here the distances of $D=85$ and $184$ Mpc to UGC~6614 and F568-6 respectively (following \\cite{pic97}; derived from a Hubble constant of $75$ km s$^{-1}$ Mpc$^{-1}$). Both galaxies show clear evidence of spiral structure (see Figures 1 and 2). Spiral arms are visible in the optical images, the \\ion{H}{1} column density maps, and as kinks in the velocity field at the level of 10--30 km/s. This spiral arm structure is coincident in the $R$ band and \\ion{H}{1} images as well as in the velocity field. The spiral structure for the two galaxies is large scale extending to a radius of more than 50 and 90 kpc in UGC~6614 and F568-6 respectively. This is in contrast to normal galaxies such as M81 where the entire spiral arm structure lies within the optical disk or within $\\sim 20$ kpc. We note however that in terms of disk scale lengths the extent of spiral structure in UGC~6614 and F568-6 is not so large, extending only to 3.6 and 5 scale lengths respectively, which is small compared to 8 scale lengths in M81 (the disk exponential scale length in M81 $\\sim 2.5$ kpc, \\cite{ken87b}). \\subsection{Spiral Arm Morphology} For both galaxies the morphology of the spiral structure consists of coherent spiral arms each of which wrap around the galaxy more than half a revolution. In this respect the spiral structure does not resemble that of flocculent late-type galaxies where pieces of arms exist only locally. Both spiral structures are strongly asymmetric, particularly in their inner regions. The outer arms in UGC~6614, however, are close to being a bi-symmetric or two arm spiral pattern which is centered about a position located to the east of the nucleus. In flocculent galaxies, the spiral structure is thought to be primarily propagating in the gas disk without strong coupling to the stellar disk (e.g. \\cite{bra93}). Also, simulations have shown that when the disk mass is low, the spiral structure is more likely to be flocculent (\\cite{sel84}, \\cite{car85}). The more coherent nature of the spiral structure observed in these low surface brightness galaxies suggests that the stellar disk is coupled to the gas disk and actively involved in the spiral structure. The \\ion{H}{1} column density map shows large arm/interarm contrasts of $\\sim$ 2:1 and velocity variations seen as kinks in the velocity field detectable at the level of 10--20 and 10--30 km/s in UGC~6614 and F568-6 respectively. We note that these values for the arm/interarm contrast and velocity residuals are similar to those observed in M81 which has radial velocity perturbations of $\\sim 10$ km/s observable in \\ion{H}{1} (\\cite{vis80}). High surface brightness galaxies can also have significantly stronger velocity perturbations. For example, M51 has radial velocity perturbations of 60--90 km/s seen in the CO velocity field (\\cite{vog88}), and UGC~2885 has perturbations of 50--70 km/s observed in H$\\alpha$ (\\cite{can93a}). Along the kinematic minor axis of the galaxy velocity residuals are expected to be radial. The pattern of the velocity residual pattern depends on whether the spiral arm pattern lies within the corotation radius (\\cite{can93}). Kinks in the velocity field alternate sign only a few times as a function of radius along the minor axes of the two galaxies. This residual pattern which is consistent with a single residual alternating sign pair associated with each arm is similar to the velocity residual field of M81 and suggests that the entire spiral pattern lies within the corotation resonance (\\cite{can93}). Low surface brightness galaxies can have gas densities (\\cite{vdh93}, \\cite{mcg92}) which fall below the threshold $\\Sigma_{crit}$ which \\cite{ken89} found was required for massive star formation in normal spiral galaxies. As proposed by \\cite{ken89} this critical density is directly related to the local stability of the disk where the Toomre stability parameter $Q$, defined as \\begin {equation} Q \\equiv {\\kappa \\sigma \\over 3.36 G \\Sigma } \\end{equation} (see \\cite{B+T}) can be written in terms of the critical gas density as \\begin {equation} Q = {\\Sigma_{crit} \\over \\Sigma \\alpha}. \\end{equation} Here $\\Sigma$ is the local gas surface density, $\\kappa$ is the epicyclic frequency, $\\sigma$ is the velocity dispersion, and $\\alpha$ was emperically determined by \\cite{ken89} to be $\\sim 0.7$. When $Q > 2$ amplification processes such as the swing amplifier are inefficient and the disk is unresponsive to tidal perturbations which could excite spiral density waves in a more unstable disk (see \\cite{B+T} and references theirin). It is therefore unlikely for a disk with $Q>2$ to show spiral structure. This implies that a disk with gas density below the critical gas density of \\cite{ken89} is also unlikely to show spiral structure assuming there is no other disk mass component. (Note that $Q\\sim 1.4 $ for $\\Sigma_{crit} /\\Sigma = 1$.) Consequently if a gas disk is well below the critical gas density and yet shows spiral structure, a natural explanation is that there is another massive component in the disk (see \\cite{jog84} for instability in a two fluid disk). Low surface brightness galaxies can show spiral structure despite the fact that their gas densities fall below the critical density (\\cite{vdh93}, \\cite{deb96b}). Unfortunately the \\ion{H}{1} data of \\cite{pic97} is not of sufficiently high velocity resolution to measure the gas velocity dispersion $\\sigma$ required to calculate the $\\Sigma_{crit}$ predicted by \\cite{ken89}. However, \\cite{pic97} found that even if $\\sigma$ had the low value of 6km/s, UGC~6614 had gas densities below the critical density for $r <50''$, $r>120''$ and that Malin 2 had gas densities below this threshold everywhere. This suggests that a gas plus stellar disk may be required for these disks to be sufficiently unstable to support the observed spiral structure. \\subsection{Spiral Arm Amplitudes} To estimate the amplitude of the spiral arm structure observed in \\ion{H}{1} and in the $R$ band images, we must first correct for the inclination of the galaxy. This is relatively straightforward in the case of UGC~6614 since the $R$ band isophote, \\ion{H}{1} isophote and \\ion{H}{1} velocity field derived position and inclination angles all agree (PA = $116^\\circ$, $i=35^\\circ\\pm 3 ^\\circ$, \\cite{pic97}). For F568-6 the $R$ band isophote and \\ion{H}{1} velocity field derived position and inclination angles agree (PA = $75 ^\\circ$, $i=38^\\circ\\pm 3 ^\\circ$), though the \\ion{H}{1} distribution is more asymmetric than the $R$ band light (\\cite{pic97}). We have adopted these orientations to correct for the galaxy inclinations. In Figures 3 and 4 we show azimuthal profiles in the inclination corrected $R$ band images and \\ion{H}{1} maps for the two galaxies. The amplitudes as a function of radius of the $m=1$ and $m=2$ azimuthal Fourier components expressed as a percentage of the azimuthal average for the two galaxies are shown in Figure 5 for the $R$ band images and in Figure 6 for the \\ion{H}{1} images. UGC~6614 has an oval distortion in its central region ($r < 40''$) which is particularly noticeable after correction for inclination. (By oval distortion we mean an elongation in the isophotes that does not vary in position angle over a range of radius.) This oval distortion is seen as a large azimuthal density variation (evident as high $m=2$ components in Figure 5) which is not an actual intensity variation in the spiral arms. In the outer regions of both galaxies the faint spiral arms are far brighter than the underlying disk which is only marginally detected. Although peaks in the \\ion{H}{1} and $R$ band surface brightness are correlated (see Figures 1 and 2) they do not correspond to maximum densities or surface brightnesses in the azimuthal cuts except at large radii in UGC~6614 (see Figures 3 and 4). Better correlation might be observed with higher angular resolution \\ion{H}{1} observations. Since the amplitudes shown in Figures 3 and 4 depend on the assumed galaxy orientations we recomputed them for moderate variations in the inclination angles. For a $5^\\circ$ lower inclination of $30^\\circ$ in UGC~6614, the change in amplitudes were largest in the region of the oval distortion ($r<40''$) and were $10-15\\%$ smaller than at an inclination of $35^\\circ$. Variations in the amplitudes elsewhere in UGC~6614 and in F568-6 for a corresponding difference in inclination were smaller, $\\lta 5\\%$. These amplitude uncertaintites are not sufficiently large to significantly change the limits for the mass-to-light ratio we estimate below. ", + "conclusions": "The spiral structure of the low surface brightness galaxies F568-6 and UGC~6614 is large scale, with arms that wrap more than half a revolution, and extend out to 50 and 80 kpc in UGC~6614 and F568-6 respectively. These spiral arms are visible in the $R$ band images, the \\ion{H}{1} column density maps and as kinks in the \\ion{H}{1} velocity fields. The density contrasts observed in the \\ion{H}{1} maps are high, with arm/interarm contrasts of $\\sim 2:1$, whereas the velocity perturbations due to spiral structure are low, in the range 10-20 km/s and 10-30 km/s in UGC~6614 and F568-6 respectively. We use the small velocity response to place upper limits on the mass-to-light ratio of the stellar disk. The strongest limits occur at small radii ($r \\sim 40''$ or 16 kpc and $r \\sim 30''$ or 24 kpc in UGC~6614 and F568-6 respectively) where there are strong distortions observed in the $R$ band images. The weak observed response in the $\\phi$ velocity component limits the mass-to-light ratios of the disk in these regions to $M/L < 3$ and 6 for UGC~6614 for F568-6 respectively. This is equivalent to requiring the densities of the disks (if they have azimuthal variations of the same size as that observed in the $R$ band images) to be less than 30 and 60 $M_\\odot/{\\rm pc}^2$ at a radius of 16 and 24 kpc for UGC~6614 and F568-6 respectively. An underlying more massive disk at these radii could exist if it had smaller azimuthal density variations than observed in the $R$ band. These limits are sufficiently strong to require a significant dark matter component even in the central regions of this galaxy, confirming the findings of previous studies. Our limits furthermore imply that this dark matter component cannot be in the form of a cold disk since a cold disk would necessarily be involved in the spiral structure, though a hot disk cannot be excluded. We note that this upper limit was derived assuming a linear gas response in the tight winding or WKB approximation, and that a non-linear gas response driven by bar like or oval distortions could cause small velocity perturbations to be present even in the case of stronger potential perturbations (or higher mass-to-light ratios). To produce the large arm/interarm \\ion{H}{1} density variations it is likely that the spiral arm potential perturbation is sufficiently strong to produce shocks in the gas. For a forcing that is greater than $2\\%$ of the axisymmetric force, $M/L \\gta 1$ is required in both galaxies in the outer regions. This is equivalent to a disk surface density between $r =$ 60--120\\arcsec\\ in UGC~6614 of 2.6--1.0 $M_\\odot/{\\rm pc}^2$ and between $r =$ 40--90\\arcsec\\ in F568-6 of 6.6--1.0 $M_\\odot/{\\rm pc}^2$ assuming that the amplitude of the variations in the disk mass is the same as that observed in the $R$ band. These lower limits imply that the stellar surface density is of the same order as the gas surface density. This is consistent with the large scale morphology of the spiral structure, and the stability of the gas disk, both of which suggest that a moderate stellar component is required to produce the observed spiral structure. The gas disks alone probably fall below the critical gas density (\\cite{pic97}) emperically found by \\cite{ken89} for the onset of masive star formation and so are not likely to be unstable enough to support spiral density waves. However it is likely that the combined stellar and gas disks are (see \\cite{jog84} for instability in a two fluid disk). The coupled gas and stellar disk would be consistent with the large scale spiral arm morphologies which do not resemble that of gas dominated flocculent galaxies. The limits for the mass-to-light ratio of the disk derived here suggest that the disks of these two low surface brightness galaxies lie in the range $1 < M/L < 6$ (in the $R$ band). This range is identical to that found for normal higher surface brightness disk galaxies derived from maximal disk fits to the observed rotation curves (\\cite{ken87a}, \\cite{ken87b}). In other words the disk mass-to-light ratio limits placed here are not abnormal compared to those of normal galaxies. Our limits on the disk surface densities remain consistent with previous studies which find that low surface brightness galaxies have substantially lower mass surface densities than normal galaxies (\\cite{deb96b}, \\cite{spray95b}). In low surface brightness galaxies since the disk contributions to the rotation curves are small compared to the halo, good fits to the rotation curves can be acheived with a range of mass-to-light ratios. Fits to the rotation curves using halo profiles of the form proposed by \\cite{nav95} yeilded the best fit mass-to-light ratios of $\\sim 0.8$ and $0.5$ for UGC~6614 and F568-6 respectively (\\cite{pic97}), whereas \\cite{imp97} using an isothermal halo found a good fit to the rotation curve of F568-6 with a much higher $M/L_B = 8$. The disk mass-to-light ratios derived from these fits probably depend upon the halo profile assumed and whether the bulge and disk are allowed to have different mass-to-light ratios and so are not tightly constrained. We note that multi-wavelength observations (particularly those in the infrared) find that amplitude variations across spiral arms can be a strong function of wavelength (e. g. \\cite{rix93}). While the low \\ion{H}{1} column depths in UGC~6614 and F568-6 suggest that extinction from dust is not a large effect in these low surface brightness galaxies, it would not be surprising if an older population of stars (which would be more apparent in the near infrared wavelengths) might have smaller spiral arm amplitudes. In this case a lower limit for the mass-to-light derived as we have done here but from a near infrared image might yield even stronger limits requiring an even more massive stellar disk. If high mass-to-light ratios are required, then a mass-to-light ratio variation across the disk could be required to yield a good fit to the rotation curve assuming a smooth halo component, and a large dark matter component might not necessarily be required in the central regions. Multi-wavelength observations coupled with metallicity measurements should also provide useful information about what kinds of stellar populations would be consistent with the range of mass-to-light ratios given here. It might prove to be interesting to place similar mass-to-light ratio limits in low surface brightness dwarf galaxies which also can have spiral structure (for example UGC~11820, UGC~5716, and F568-1, from \\cite{vz97} and \\cite{deb96b}). These galaxies can have regular, symmetrical \\ion{H}{1} velocity fields and smooth \\ion{H}{1} column density maps, even when the optical components lack symmetry or contain strong spiral structure. The optical components therefore don't strongly influence the velocity field (although some of these effects might be visible in \\ion{H}{1} observations at higher angular resolution). This suggests that the upper limits discussed here might be particularly revealing. Strong spiral structure detected in \\ion{H}{1} in the outer regions of higher surface brightness galaxies could be used to estimate the mass density of a low surface brightness stellar disk. Some interesting candidates for such a study might be the dark blue compact dwarf NGC~2915 which may contain spiral structure well outside its optical disk (\\cite{meu96}), and the polar ring galaxy NGC~4650A which shows evidence for spiral arms in its ring (\\cite{arn97})." + }, + "9701/astro-ph9701190_arXiv.txt": { + "abstract": "We present radio observations of the black hole candidate X-ray binary GX 339-4 with the Australia Telescope compact array. Mapping of the highest resolution 3.5 cm data reveals a jet-like extension, which if confirmed would be the first detection of a radio jet from a {\\em persistent} black-hole candidate system. No evidence is found for associated structures such as bow shocks or jet lobes on larger scales. The spectral energy distribution from 22 - 3 cm is relatively flat, suggesting emission is dominated by a compact absorbed core. ", + "introduction": "Radio emission is an observational property of $\\sim 15$\\% of X-ray binaries (see e.g. Hjellming \\& Han 1995). It is anticorrelated with the property of X-ray pulsations (Fender et al. 1997) and displays the nonthermal spectra and high (T$_b > 10^8$ K) brightness temperatures characteristic of synchrotron emission. Radio jets from X-ray binaries were first discovered in the late 1970s from SS 433 (Spencer 1979) in which the ejecta were subsequently observed to have an ejection velocity of 0.26 c (Hjellming \\& Johnston 1981). Over the ensuing $\\sim 15$ years radio-jets were discovered from perhaps six further X-ray binary sources, in none of which however accurate proper motions of individual ejecta, or plasmons, could be measured (see e.g. Fender, Bell Burnell \\& Waltman 1997 for a review of observational properties). In the past three years however, apparent superluminal plasmon velocities, corresponding to true relativistic bulk motions at $\\sim 0.9$c have been observed from two transient X-ray binaries, GRS 1915+105 (Mirabel \\& Rodriguez 1994) and GRO J1655-40 (Tingay et al. 1995). GX 339-4 is an unusual X-ray binary which exhibits distinct and varied X-ray states, but is in all states persistently and clearly observable with X-ray telescopes, unlike the transient radio-jet sources GRS 1915+105 and GRO J1655-40. The characteristics of the X-ray emission, notably a hard tail, rapid variability and a similarity to Cyg X-1, have lead to classification of the source as a strong black hole candidate (e.g. Makishima et al. 1986; Miyamoto et al. 1992; Miyamoto et al. 1993; Nowak 1995). Optical studies (Callanan et al. 1992) reveal a 14.8 hr modulation, probably orbital in origin, in the light from the 16-19th magnitude variable optical counterpart. Distance estimates to the source range from 1.3 kpc (Predehl et al. 1991) to $\\geq 4$ kpc (Makishima et al. 1986). The discovery of radio emission from the source was announced in Sood \\& Campbell-Wilson (1994), and sporadic monitoring has been ongoing over the past $\\sim 18$ months at 36 cm using the Molonglo Observatory Synthesis Telescope (MOST). The source was also detected during the southern hemisphere radio survey of X-ray binaries conducted by Spencer et al. (1997). Durouchoux et al. (1997) and Sood et al. (1997) discuss the radio behaviour of the source over the past $\\sim 18$ months in the context of X-ray monitoring and similarities to Cyg X-1. ", + "conclusions": "GX 339-4 has become the latest black hole candidate X-ray binary to be detected as a variable radio source. Durouchoux et al. (1997) and Sood et al. (1997) have already discussed possible radio - X-ray correlations in the source, and have compared it to Cygnus X-1, the most famous persistent black hole candidate radio source. In four days of radio observations with ATCA in high-resolution configuration we have measured the spectral energy distribution of GX 339-4 from 22 - 3 cm and mapped the region of the source on scales from arcmin to arcseconds. We find a relatively flat ($\\alpha \\sim +0.2$) radio spectrum for the emission, indicative of a compact absorbed core. Low-resolution mapping finds no evidence on arcmin-scales for lobes. bow shocks or synchrotron nebulae such as those associated with other radio-jet X-ray binaries. Most intriguingly however we have presented evidence for the existence of a jet-like structure in radio maps of GX 339-4. If confirmed, this source becomes the first persistent black-hole candidate source to reveal a radio jet. The importance of the detection of a jet, which is in all likelihood relativistic, from a non-transient black hole candidate X-ray binary cannot be understated. Both GRS 1915+105 and GRS J1655-40 are transient sources, and while displaying quasi-recurrent behaviour in the 2-3 years since their discovery, the lack of detection of these systems by previous X-ray satellites (Lochner \\& Whitlock 1992; Castro-Tirado 1994; Castro-Tirado, private communication) testifies to prolonged states below detectable levels. GX 339-4 on the other hand is a persistently visible source, both at X-rays and (at least for past $\\sim 18$ months) in the radio, and may turn out to be one of our most valuable laboratories in the study of radio jets." + }, + "9701/hep-ph9701416_arXiv.txt": { + "abstract": "We simulate the formation of cosmic strings at the zeros of a complex Gaussian field with a power spectrum $P(k) \\propto k^n$, specifically addressing the issue of the fraction of length in infinite strings. We make two improvements over previous simulations: we include a non-zero random background field in our box to simulate the effect of long-wavelength modes, and we examine the effects of smoothing the field on small scales. The inclusion of the background field significantly reduces the fraction of length in infinite strings for $n < -2$. Our results are consistent with the possibility that infinite strings disappear at some $n = n_c$ in the range $-3 \\le n_c < -2.2$, although we cannot rule out $n_c = -3$, in which case infinite strings would disappear only at the point where the mean string density goes to zero. We present an analytic argument which suggests the latter case. Smoothing on small scales eliminates closed loops on the order of the lattice cell size and leads to a ``lattice-free\" estimate of the infinite string fraction. As expected, this fraction depends on the type of window function used for smoothing. ", + "introduction": "Cosmic strings are effectively one-dimensional topological defects which may form at a phase transition in the early universe (see Ref. \\cite{review} for a review). Although much of the early interest in cosmic strings has centered on the possibility that they might have served as seeds for the formation of large-scale structure, cosmic strings are interesting physical objects in any case, and they have analogues in the study of condensed matter \\cite{condmat}. Any investigation of the evolution and cosmological consequences of cosmic strings must begin with the study of the initial cosmic string configuration, a study which was first undertaken by Vachaspati and Vilenkin \\cite{VV}, and re-examined by many others (\\cite{SF} - \\cite{RY}). Although subsequent cosmic string evolution will erase many of the details of the initial configuration, one fundamental property of the initial conditions is crucial to the subsequent evolution: the existence of infinite strings. Without the existence of infinite strings, the cosmic strings produced at the phase transition may all decay via gravitational radiation long before they can have any interesting cosmological effects. Vachaspati and Vilenkin found in their simulation that roughly 80\\% of the string was in the form of infinite strings \\cite{VV}. Using a very different type of simulation for the string formation process, Borrill \\cite{Borrill} claimed that this fraction was, in fact, zero, while Robinson and Yates \\cite{RY}, in a study of the dependence of this fraction on the power spectrum of the initial field, argued that for power spectra of the form $P(k) \\propto k^n$, the infinite string fraction $f_\\infty$ drops to zero for $n \\le -2$. In fact, this result is not obvious from their simulations; it is based on fitting $f_\\infty$ to an analytic function for $n > -2$. In this paper, we extend the simulations of Robinson and Yates in two ways. First, we include the effects of long-wavelength modes which are absent from earlier simulations. Second, we examine the effect of smoothing the initial field to remove lattice effects. We find that the former change has a dramatic effect on $f_\\infty$, sharply reducing $f_\\infty$ for $n < -2$. Smoothing also affects $f_\\infty$, leading to ``lattice-free\" estimates of the fraction in infinite strings. In the next section we present our numerical results, and in Section 3 we discuss briefly our main conclusions. ", + "conclusions": "Our results are consistent with the possibility that a transition occurs in the string network at some value $n= n_c$, with $-3 \\le n_c < -2.2$; when $n < n_c$ the infinite strings disappear from the network, leaving only closed loops. Robinson and Yates \\cite{RY} argued that $n_c=-2$, but this value is not consistent with our results. Furthermore, we cannot rule out the possibility that $n_c = -3$, i.e., the infinite strings do not disappear until the mean string density goes to zero. In fact, the following argument suggests that infinite strings should not disappear for $n > -3$. The real and imaginary parts of $\\phi$ are independent Gaussian fields, and the zeros of each of these fields form a set of two-dimensional surfaces. Consider first the surfaces defined by Re$(\\phi) = 0$. Since the volumes of space occupied by the regions with positive and negative Re($\\phi$) are equal, we expect both regions to percolate to an infinite distance. Hence, the boundaries dividing these regions should contain at least some infinite surfaces. This argument holds for both the surfaces defined by Re($\\phi) = 0$ and Im($\\phi) = 0$. The intersection of these two sets of surfaces gives us the location of the cosmic strings. Since infinite surfaces must exist in both sets of surfaces, it appears that there must also be infinite cosmic strings. Note that it is possible to imagine rather arcane distributions of the fields which violate this argument. For example, the regions with Re$(\\phi) > 0$ and Re$(\\phi) < 0$ could be nested inside of each other in larger and large finite volumes, producing a fractal distribution with arbitrarily large but finite surfaces of Re$(\\phi) = 0$. However, it seems unlikely that a Gaussian field could lead to such a distribution. The crucial point in this argument is the fact that the distribution is symmetric with respect to positive and negative values of Re($\\phi$) and Im($\\phi$); if we relax this assumption and ``bias\" the distribution, the argument no longer holds. In fact, infinite strings are observed to disappear in simulations with such a ``bias\" \\cite{vachaspati}, \\cite{kibble}, \\cite{HS}. The aim of our simulations with a smoothed field $\\phi$ was to obtain a lattice-independent estimate of $f_{\\infty}$. We considered only the white noise spectrum, $n=0$, and found $f_{\\infty}\\approx 0.7$ for spherical Gaussian smoothing and $f_{\\infty}\\approx 0.8$ for a sharp cut-off in $k$-space. These values are comparable to those obtained in earlier lattice simulations without smoothing \\cite{VV,SF}. The variation of $f_{\\infty}$ for different choices of smoothing is not surprising, since, for example, the total string density $L/V$ in equations (\\ref{L/V}) and (\\ref{k2}) clearly depends on the short-wavelength behavior of the spectrum $P(k)$ (which is affected by the smoothing). This variation is again comparable to the variation between the values of $f_{\\infty}$ in simulations without smoothing on different types of lattice \\cite{VV}, \\cite{SF}, \\cite{HS}. We find no evidence supporting the hypothesis \\cite{Borrill} that the presence of infinite strings is due entirely to the lattice and that they should disappear in lattice-free simulations. The smoothing length we used was sufficiently large to ensure that the loops making the largest contribution to the total string length had sizes much greater than the lattice cutoff, so that the lattice effects were minimal. Still, we found a substantial fraction of the total length to be in infinite strings. We are currently investigating the formation of domain walls and monopoles with correlated fields. Given that a parallel literature on this subject exists in condensed matter physics \\cite{bradley}, these results may have applications beyond the purely cosmological. It is conceivable, for example, that the decline of the infinite string density with the decrease of the spectral index $n$ can be tested experimentally. What one needs is a condensed matter system (such as liquid He$^4$) in which linear defects are formed at a second-order phase transition. Defect formation can then be observed by a rapid temperature (or pressure) quench from above to below the transition point \\cite{Zurek,Hendry}. Near the critical temperature $T_c$, the order parameter develops long-range fluctuations. At $T=T_c$, the fluctuation spectrum is a power law of the form (\\ref{power}) with $n=-2+\\eta$, where the critical exponent $\\eta$ is typically a small number, $\\eta\\lesssim 0.05$ \\cite{Huang} ($\\eta\\approx 0.05$ for He$^4$). If the system is allowed to equilibrate very close to the critical point and is then rapidly quenched to subcritical temperatures, one can expect the length in infinite strings to be suppressed compared to a quench from a temperature well above $T_c$ (where the fluctuation spectrum is close to $n = 0$). Our Fig. 2 suggests a suppression roughly by an order of magnitude, while the value of $n_c=-2$ conjectured by Robertson and Yates \\cite{RY} would give a much more dramatic suppression. It should be noted that a realistic quench is a rather complicated process, and its outcome can depend on a variety of physical effects (for a recent discussion see \\cite{Zurek,K})." + }, + "9701/astro-ph9701123_arXiv.txt": { + "abstract": "Results of an ASCA X-ray observation of NGC 4507 are presented. The spectrum is best parameterized by a double power law (or a partial covering) model plus a narrow FeK$\\alpha$ emission line likely due to the line-of-sight absorption matter. The data require also an emission line at $\\sim$ 0.9 keV consistent with NeIX, which may indicate that the soft X-ray emission derives from a combination of resonant scattering and fluorescence of the radiation by photoionized gas. Moreover, complex absorption (or other, unresolved, emission lines) is required by the data below 3 keV. Similarities between the X-ray spectra of NGC 4507 and NGC 4151 are also stressed. ", + "introduction": " ", + "conclusions": "" + }, + "9701/astro-ph9701208_arXiv.txt": { + "abstract": " ", + "introduction": "Topological defect theories provide an alternative to inflation for explaining the origin of the primordial fluctuations which have grown into present-day structures through gravitational instability. Cosmic strings are one-dimensional topological defects possibly formed in a phase transition in the early Universe. The observed presence of massive nonlinear structures at high redshifts of 3 to 4 has provided stringent constraints on models of structure formation containing hot dark matter (HDM), a candidate for which is a massive neutrino. This is due to the large thermal velocities of the HDM particles, which prevent the growth of perturbations on scales smaller than the free-streaming length. For inflationary models with adiabatic density fluctuations, recent data on the abundance of damped Lyman alpha absorption systems (DLAS)$^{\\cite{DLAS}}$ and on the quasar abundance$^{\\cite{QSO}}$ has restricted the fraction of HDM to less than about $30\\%$ of the total dark matter present$^{\\cite{MaB}}$. The scenario of structure formation with cosmic strings is, however, viable even if the dark matter in the universe is hot, because cosmic strings, which provide the seeds for the density perturbations, survive the neutrino free-streaming$^{\\cite{VS 83,BKST}}$. In this article we present work on the constraints imposed on the cosmic string theory by the abundance of high redshift quasars$^{\\cite{MB}}$. Searches for high redshift quasars have been going on for some time$^{\\cite{QSO}}$. The quasar luminosity function is observed to rise sharply as a function of redshift $z$ until $z \\simeq 2.5$. According to recent results from the Palomar grism survey by Schmidt et al. (1995)$^{\\cite{QSO}}$, it peaks in the redshift interval $z \\, \\epsilon$ [1.7, 2.7] and declines at higher redshifts. Quasars (QSO) are extremely luminous, and it is generally assumed that they are powered by accretion onto black holes. It is possible to estimate the mass of the host galaxy of the quasar as a function of its luminosity, assuming that the quasar luminosity corresponds to the Eddington luminosity of the black hole. For a quasar of absolute blue magnitude $M_B= -26$ and lifetime of $t_Q=10^8 yrs$, the host galaxy mass can be estimated as$^{\\cite{PS 94}}$ \\be M_G=c_1 10^{12} M_\\odot \\, , \\ee where $c_1$ is a constant which contains the uncertainties in relating blue magnitude to bolometric magnitude of quasars, in the baryon fraction of the Universe and in the fraction of baryons in the host galaxy able to form the compact central object (taken to be 10$^{-2}$). The best estimate for $c_1$ is about 1. Models of structure formation have to pass the test of producing enough early objects of sufficiently large mass to host the observed quasars. Besides the quasars themselves, absorption lines in their spectra due to intervening matter can be used to quantify the amount of matter in nonlinear structures at high redshifts. Based on the number density of absorption lines per frequency interval and on the column density calculated from individual absorption lines, the fraction of $\\Omega$ in bound neutral gas (denoted by $\\Omega_g$) can be estimated. For high-column density systems, the damped Ly-$\\alpha$ systems (DLAS), recent observational results are that \\be \\Omega_g (z) > 10^{-3} \\ee for $z \\, \\epsilon$ [1,3]. In the most recent results by Storrie-Lombardi et al. (1996)$^{\\cite{DLAS}}$ there is evidence for a flattening of $\\Omega_g$ at $z \\sim 2$ and a possible turnover at $z \\sim 3$ . The corresponding value for $\\Omega$ in bound matter is larger by a factor of $f^{-1}_b$, where $f_b$ is the fraction of bound matter which is baryonic, which is about $10\\%$ in a flat universe. In the next section, we give a brief review of the cosmic string scenario of structure formation. Then we study the accretion of hot dark matter onto cosmic string loops, which seed large amplitude local density contrasts already at early times. We use the results to compute the number density of high redshift objects as a function of a parameter $\\nu$ which determines the number density of loops in the scaling solution (see Section 2). We demonstrate that for realistic values of $\\nu$, the number of massive nonlinear objects at redshifts $\\le 4$ satisfies the recent observational constraint of quasar abundances (see Figure~1). We also comment on the implications of (2). We consider a spatially flat Universe containing HDM and baryons. Units where $c=1$, a Hubble constant of $H=50 \\,h_{50} {\\rm km s^{-1} Mpc^{-1}}$ and a redshift at equal matter and radiation of $z_{eq}=5750 \\,\\Omega h_{50}^{2}$ are used. ", + "conclusions": "We have studied the accretion of hot dark matter onto moving cosmic string loops and made use of the results to study early structure formation in the cosmic string plus HDM model. The loop accretion mechanism is able to generate nonlinear objects which could serve as the hosts of high redshift quasars much earlier than the time cosmic string wakes start becoming nonlinear (which for $G \\mu = 10^{-6}$ and $v_s = 1/2$ occurs at a redshift of about 1). The fraction $\\Omega_{nl} (z)$ of the total mass accreted into nonlinear objects by string loops unfortunately depends very sensitively on $\\alpha$ and $\\nu$. On the other hand, this is not surprising since the power of the loop accretion mechanism depends on the number and initial sizes of the loops, and the scaling relation $\\Omega_{nl} \\sim \\nu \\alpha$ is what should be expected from physical considerations. For the values $\\nu =1$ and $\\alpha_{-2} = 1$ which are indicated by recent cosmic string evolution simulations$^{\\cite{CSsim}}$, we conclude that the loop accretion mechanism produces enough large mass protogalaxies to explain the observed abundance of $z \\leq 4$ quasars (see Figure~1). Note that the amplitude of the predicted protogalaxy density curves depends sensitively on the parameters of the cosmic string scaling solution which are still poorly determined. Hence, the important result is that there {\\it are} parameters for which the theory predicts a sufficient number of protogalaxies. Since not all protogalaxies will actually host quasars, and since the string parameters are still uncertain, it would be wrong to demand that the amplitude of the $n_G$ curve agree with that of the observed $n_Q$; rather, it should lie above the $n_Q$ curve. It is more difficult to make definite conclusions regarding the abundance of damped Lyman alpha absorption systems. In the form of Eq.~\\ref{Onlres}, the condition for the cosmic string loop accretion mechanism to be able to explain the data is also satisfied. However, Eq.~\\ref{Onlres} refers to the value of $\\Omega$ in baryonic matter. The corresponding constraint on the total matter collapsed in structures associated with damped Lyman alpha systems is $$ \\Omega_{\\rm DL} (z < 3) > f^{-1}_b 10^{-3} $$ where $f_b$ is the local fraction of the mass in baryons. From Eq. 5 it follows that the above constraint is only marginally satisfied, and this only if the local baryon fraction $f_b$ exceeds the average value for the whole Universe of about $f_b = 0.1$. But in the cosmic string model with HDM we might expect $f_b$ in nonlinear objects to be enhanced over the average $f_b$ because after $t_{eq}$ baryons are able to cluster during the time that the HDM is prevented from accreting by the free streaming. Thus, cosmic strings may be able to restore agreement with (~\\ref{Onlres}) in a natural way. More calculations are required to resolve this issue. Here we have only reported on the mechanism of forming early nonlinear objects through accretion onto string loops. Another mechanism is through small-scale structure on the long strings leading to the formation of filaments rather than wakes, which has recently been investigated$^{\\cite{ZLB}}$. It was found that this could be the most effective mechanism, and for the maximal possible amount of small-scale structure, $\\Omega_{nl} \\sim 1$ can be reached already at a redshift of $5$. It is clearly important to determine the amount of small-scale structure present on strings." + }, + "9701/astro-ph9701178_arXiv.txt": { + "abstract": "We discuss general relativistic effects in the steady-state neutrino-driven ``wind'' which may arise from nascent neutron stars. In particular, we generalize previous analytic estimates of the entropy per baryon $S$, the mass outflow rate $\\dot M$, and the dynamical expansion time scale $\\tau_{\\rm dyn}$. We show that $S$ increases and $\\tau_{\\rm dyn}$ decreases with increasing values of the mass-to-radius ratio describing the supernova core. Both of these trends indicate that a more compact core will lead to a higher number of neutrons per iron peak seed nucleus. Such an enhancement in the neutron-to-seed ratio may be required for successful $r$-process nucleosynthesis in neutrino-heated supernova ejecta. ", + "introduction": " ", + "conclusions": "" + }, + "9701/astro-ph9701102_arXiv.txt": { + "abstract": "The dust extinction towards bright \\hrs in NGC 598 and NGC 5457 has been studied in detail by forming line ratios of Balmer and Paschen emission lines covering a large wavelength range. Three homogeneous models of the geometrical distribution of the emitting sources and the obscuring dust have been tested. Only for low extinctions can the data be fit by a homogeneous slab. For most of the observed \\hrs~ the Witt et al. (\\cite{WTC}) 'dusty nucleus' model matches the observations equally well as the usual assumption of a foreground screen, but the former implies much larger actual dust contents. Spatial variations in the optical depth in V of the order 0.3 across a region are found. ", + "introduction": "Investigations of external galaxies are often faced with the problem of disentangling the effects of dust and the physical properties, e.g. age, IMF, chemical composition etc. of the emitting source. The effect of dust extinction goes roughly as $\\lambda^{-1}$ and will consequently change the observed spectral energy distribution. Studies and interpretation of the physical properties in external galaxies therefore rely on how well we can account for the attenuation of the emitted light due to dust. The wavelength dependence makes multiwavelength analyses favourable and a growing number of investigations combining optical and IR wavelengths have been made both using imaging (e.g. Jansen et al. \\cite{jan}; Witt et al. \\cite{wi94}; Evans \\cite{ev}; van Driel et al. \\cite{driel}) and spectroscopy (e.g. Puxley \\& Brand \\cite{pu:br}; Calzetti et al. \\cite{CKS96}). In the analysis of spectra of extragalactic \\hrs~ the correction for dust extinction is usually made on the assumption that the intrinsic hydrogen emission line spectrum is known through case B recombination and any difference between the emergent and intrinsic flux ratios can be ascribed to dust obscuration following the Galactic extinction curve. Throughout the years, the use of the ratio of the two strongest Balmer lines, H$\\alpha$/H$\\beta$, sometimes in combination with H$\\gamma$, has been prevailing, albeit they only span a limited wavelength range, because they will be included in almost any optical spectra and are easy to detect. It has long been proposed to expand the baseline for extinction determination by comparing near-IR hydrogen Paschen lines with emission lines from the Balmer series at short wavelengths (cf. Greve et al. \\cite{gr89}; Osterbrock \\cite{os89}). Such line pairs are all separated by wide wavelength intervals over which the extinction has a large effect and is therefore potentially easier to deduce. Additionally, some line pairs form corresponding multiplet lines, ${\\rm P}n / {\\rm H}n$ originating at the same upper atomic level with relative strength depending primarily on the transition probability and minimised dependence on theoretical recombination line calculations (Greve et al. \\cite{gr94}). In Petersen \\& Gammelgaard (\\cite{pe:ga}; hereafter Paper I) we have demonstrated the feasibility of observing lines from the Paschen series and the short end of the Balmer series from giant extragalactic \\hrs~ simultaneously with a two-channel spectrograph. In the present paper we will further discuss dust extinction towards \\hrs~ in NGC 598 and NGC 5457 as derived from spectra obtained with the same instrument. The straightforward application of the Galactic interstellar extinction curve derived from point sources to extended emission regions in external galaxies is, however, problematic, because this procedure implicitly assumes a dust configuration of a homogeneous foreground absorbing screen. Intuitively one would expect things to be more complicated in reality with dust distributed in a patchy or clumpy way maybe mixed with the emitting source (Witt et al. \\cite{WTC}; hereafter WTC; Calzetti et al. \\cite{cal}, \\cite{CKS96}; Puxley \\& Brand \\cite{pu:br}) and with scattering of, particularly blue, photons into the line of sight by dust grains (Bruzual et al. \\cite{bru}; WTC). This could all lead to a much greyer effective extinction curve than the standard Galactic curve. This was exemplified in Paper I by using an extinction curve with $R_{\\rm V} = A_{\\rm V} / E_{\\rm B - V} = 6$, here we will address this issue by fitting three different models of relative spatial distribution of emitting source and dust to the measured line ratios. ", + "conclusions": "The observations of giant \\hrs~ in the spiral galaxies NGC 598 of NGC 5457 have demonstrated how the dust extinction can be derived on the basis of multiple emission lines from the Balmer and Paschen series in the optical and near-IR regimes obtained simultaneously with ordinary CCD detectors. Although some of these lines can only be observed with larger flux errors the wide wavelength interval spanned by the lines and the utilisation of the redundancy of the Paschen lines at $\\lambda\\lambda$8500--9000 \\AA~ leads to a small overall uncertainty compared to the use of the two strongest Balmer lines. Inclusion of lines beyond 1 $\\mu$m offers a chance to make simple test of the geometrical distribution of emitting source and dust in the light of the common but questionable assumption of a foreground screen. A slab of homogeneous mixture can only match the data points for the \\hrs~ with relative low dust extinction. Generally we can not discriminate between the simple foreground extinction screen and the WTC 'dusty nucleus' model, although signs of deviation from the pure foreground screen are seen at near-IR wavelengths. This is in agreement with Calzetti et al. (\\cite{CKS96}), who found that the reddening towards starburst regions in 13 galaxies by large can be explained by foreground dust. The choice of extinction model for reddening correction of diagnostic line ratios only has small effects on the derived physical parameters when lines are not widely separated. It is, however, clearly seen how the foreground screen always yields a lower limit to the dust content and that considerable larger amounts of dust can be present in the 'dusty nucleus' configuration. A more definitive discrimination between various geometries, possibly including patchiness, and assessment of the actual dust content will need investigations at even longer wavelengths in combination with optical lines (Puxley \\& Brand \\cite{pu:br}; Genzel et al. \\cite{gen95}; Calzetti et al.\\cite{CKS96}) and more refined modelling (Witt \\& Gordon \\cite{wi96}). While a lot can be learned from high-S/N observations of P$\\gamma$ obtainable with the high quantum efficiency of the new generation of CCD's, it is desirable to have data on emission lines further out in the near-IR, e.g. P$\\beta$, Br$\\gamma$, Br$\\alpha$, which can not be observed with the same instruments as the blue optical lines. The kind of studies presented in this paper can serve to bridge the two wavelength ranges. One obvious result from the present study is the demonstration of a spatial variation in the optical depth towards the emission regions of the order 0.2--0.3 on small angular scales. More detailed investigations of one of the \\hrs~ suggest that the extinction has it highest value at or close to the edge of the region. This can explain at least some of the disagreement in the quoted extinction of \\hrs~ by different authors. It also emphasizes the importance of sampling exactly the same slit position and aperture size if combining data from different observations, and the care that most be taken when comparing physical parameters from miscellaneous sources in the literature." + }, + "9701/astro-ph9701044_arXiv.txt": { + "abstract": "We have carried out high resolution MHD simulations of the nonlinear evolution of Kelvin-Helmholtz unstable flows in $2\\frac{1}{2}$ dimensions. The modeled flows and fields were initially uniform except for a thin shear layer with a hyperbolic tangent velocity profile and a small, normal mode perturbation. These simulations extend work by \\cite{fran96} and \\cite{maletal96}. They consider periodic sections of flows containing magnetic fields parallel to the shear layer, but projecting over a full range of angles with respect to the flow vectors. They are intended as preparation for fully $3D$ calculations and to address two specific questions raised in earlier work: 1) What role, if any, does the orientation of the field play in nonlinear evolution of the MHD Kelvin-Helmholtz instability in $2\\frac{1}{2}$D. 2) Given that the field is too weak to stabilize against a linear perturbation of the flow, how does the nonlinear evolution of the instability depend on strength of the field. The magnetic field component in the third direction contributes only through minor pressure contributions, so the flows are essentially $2$D. In \\cite{fran96} we found that fields too weak to stabilize a linear perturbation may still be able to alter fundamentally the flow so that it evolves from the classical ``Cat's Eye'' vortex expected in gasdynamics into a marginally stable, broad laminar shear layer. In that process the magnetic field plays the role of a catalyst, briefly storing energy and then returning it to the plasma during reconnection events that lead to dynamical alignment between magnetic field and flow vectors. In our new work we identify another transformation in the flow evolution for fields below a critical strength. That we found to be $\\sim 10$\\% of the critical field needed for linear stabilization in the cases we studied. In this ``very weak field'' regime, the role of the magnetic field is to enhance the rate of energy dissipation within and around the Cat's Eye vortex, not to disrupt it. The presence of even a very weak field can add substantially to the rate at which flow kinetic energy is dissipated. In all of the cases we studied magnetic field amplification by stretching in the vortex is limited by tearing mode, ``fast'' reconnection events that isolate and then destroy magnetic flux islands within the vortex and relax the fields outside the vortex. If the magnetic tension developed prior to reconnection is comparable to Reynolds stresses in the flow, that flow is reorganized during reconnection. Otherwise, the primary influence on the plasma is generation of entropy. The effective expulsion of flux from the vortex is very similar to that shown by \\cite{weiss66} for passive fields in idealized vortices with large magnetic Reynolds numbers. We demonstrated that this expulsion cannot be interpreted as a direct consequence of steady, resistive diffusion, but must be seen as a consequence of unsteady fast reconnection. ", + "introduction": "Weak magnetic fields threading conducting fluid media can play vital dynamical roles, even when traditional criteria, such as relative magnetic and gas pressures, suggest the fields are entirely negligible. Perhaps the best example of this is the destabilizing influence of a vanishingly small magnetic field crossing a Keplerian accretion disk (\\cite{balhal91}), where the mere presence of the field seems fundamentally to alter the local flow properties. Other examples abound, however, that could be particularly important in astrophysics. Among them, we would include weak fields penetrating turbulent or otherwise strongly unstable flows, where such fields significantly alter evolution and transport properties, (\\eg \\cite{biswel89}; \\cite{catvain91}; \\cite{noret92}; \\cite{junetal95}). Sheared motion is a critical common element in many of these flows, and the consequent stretching of a weak, but large-scale field can lead to a locally enhanced role for the field. The flows will also frequently lead to current sheets and associated magnetic field topologies unstable to reconnection, and that is central to the nonlinear evolution of the systems (\\eg \\cite{bis93};~\\cite{park94}). Through these processes the fields can also have more global consequences. Study of the nonlinear evolution of the classical Kelvin-Helmholtz (KH) instability could be particularly useful as a well-defined example of strongly sheared flows. Further, since KH unstable boundary layers are probably common, the behavior of the instability is important for its own sake. Although the KH instability is fairly well studied in ordinary hydrodynamics (\\eg \\cite{corsher84}), comparable study has been much slower in magnetohydrodynamics (MHD). That is because the magnetic field substantially complicates the physics itself and also because computational methods and resources needed for such studies are only recently up to the task. The linear analysis of the MHD KH instability is relatively straightforward and was long ago carried out for a number of simple flow and field configurations (\\eg \\cite{chan61}; \\cite{miupr82}). Generally, and especially if the velocity change is not supersonic, the ordinary fluid shear layer is unstable to perturbations with wave vectors in the plane of the shear layer and with wavelengths greater than the thickness of the layer (\\eg \\cite{miura90}). When there is a field component projecting onto the flow field, magnetic tension provides a stabilizing influence. A simple vortex sheet is stabilized against linear perturbations whenever the magnetic field strength is sufficient that $c_a > |(\\hat {\\bmit k} \\cdot {\\bmit U_o})/ (2 \\hat {\\bmit k} \\cdot \\hat{\\bmit B_o})|$, where $\\bmit U_o$ is the velocity difference between the two layers, $c_a$ is the Alfv\\'en speed, ${\\bmit k}$ is the perturbation wave vector, and $\\hat{\\bmit B_o}$ is the direction of the magnetic field (\\cite{chan61}). \\cite{fran96} (Paper I) and \\cite{maletal96} (MBR) recently presented complementary nonlinear analyses of the MHD KH instability in mildly compressible flows based on two-dimensional numerical simulations carried out with new (and different) Riemann-solver-based MHD codes. While not the first numerical studies of the MHD KH instability, they represented big improvements over previous calculations in both numerical resolution and extent to which flow evolution was followed towards asymptotic states (readers are referred to Paper I for additional, earlier citations). Considering perturbed ``2D'' flows that were uniform except for a thin, smooth velocity transition layer, those two papers emphasized the qualitatively different behaviors in the nonlinear evolution of unstable flows depending on how close the field strength is to its critical strength for stabilization. For fields only slightly below the critical value, enhancements in the tension of the field through linear growth can stabilize the flow before it develops distinctly nonlinear characteristics. For weaker fields, however, the initial evolution of the instability is very similar to that for the ordinary KH instability. That results in the formation of eddies, and hence to substantial stretching of the magnetic field lines as well as reconnection. Paper I emphasized the remarkable fact that in a case with a field 2.5 times weaker than critical, reconnection can lead to self-organization in the flow and fairly rapid relaxation to a quasi-steady laminar and marginally stable flow. MBR presented summaries of simulations extending to somewhat weaker fields showing evidence for similar behaviors. Neither Paper I nor MBR, however, explored the problem in sufficient depth to establish the conditions necessary for the previously mentioned self-organization. In addition, it is very important in this situation to understand how the magnetic fields behave when they are ``very'' weak (a concept whose definition needs clarification, in fact). A closely related matter is what differences, if any, exist between the behavior of a truly weak field and a stronger field whose projection onto the flow vectors is weak. Alternatively stated, are there differences between the nonlinear ``2D'' MHD KH instability and the ``$2\\frac{1}{2}$D'' MHD KH instability? Answers to those basic questions are the objective of this paper. We find: 1) for the cases we have considered with an initially uniform field that the magnetic field transverse to the plane is unimportant and 2) there is a transition from the role of the magnetic field as a catalyst to flow self-organization to a role as an added source of energy dissipation that should vanish directly as the initial magnetic field strength projected onto the plane vanishes. Ultimately we must understand the full ``3D'' problem, in which the perturbation wave vector also lies outside the flow direction. On the other hand it has been difficult to carry out 3D MHD simulations with sufficient numerical resolution to be confident of the results in complex flows such as these. In addition, it will be useful to compare fully 3D behaviors with 2D flows. We hope that the current work is a significant, constructive step towards a full understanding of this problem. The paper plan is as follows. In \\S 2 we will summarize the problem set-up and relevant results from Paper I. \\S 3 contains a discussion of new results, while \\S 4 provides a brief summary and conclusion. We also include an appendix presenting an analytical model for diffusive flux expulsion from a steady vortex, in order to contrast that physics with what we observe in the eddies that form in our simulations. ", + "conclusions": "We have carried out a series of high resolution MHD simulations of Kelvin-Helmholtz unstable flows in $2\\frac{1}{2}$ dimensions. All of these simulations involve magnetic fields initially too weak to stabilize the flows in the linear regime; \\ie $B_{p0} < B_c$. Thus, since simulations are performed on a periodic space, flows all begin formation of a single ``cat's eye'' vortex. If the field lying in the computational plane is absent or ``very weak'' the cat's eye structure becomes a persistent, stable feature that represents a ``quasi-steady'' equilibrium. When there are ``very weak'' magnetic fields in the plane they become wrapped into the vortex and amplified by stretching. However, within a single turn of the vortex they are subject to tearing mode instabilities leading to magnetic reconnection. That reconnection isolates some magnetic flux within the vortex, which is eventually annihilated. This is the process through which flux is effectively expelled from a vortex. As long as the vortex persists this process will repeat. Since reconnection is irreversible, this process is also dissipative and leads to an increase over viscous effects in conversion from kinetic to thermal energy. We find in this regime that as the initial magnetic field within the computational plane is increased the dissipation rate increases in a similar manner. Likewise, as we use a finer numerical grid, thus {\\it reducing the effective numerical resistivity and viscosity, the dissipation rate increases}, reflecting the increased ability of our code to capture small-scale reconnection events. This trend is backwards from what one would expect if simple magnetic diffusion were primarily responsible for the reconnection. It suggests, perhaps that if we had been able to extend these calculations to even higher resolution the energy dissipation rate might have converged to a value independent of the effective resistivity, just as some studies of resistive MHD turbulence find. The reconnection and expulsion of flux within vortices in our simulations are similar to those in a classic study by Weiss of vortex flux expulsion in large-magnetic-Reynolds-number flows. If the initial magnetic field is strong enough that within a single turn of the vortex it is amplified around the vortex perimeter to ``dynamical'' strength ($B^2 \\sim \\rho U^2$), then the reconnection described in the previous paragraph releases stresses that are capable of disrupting the vortex entirely. This can happen in a single event or, if the field is only marginally strong enough ($B_{p0} \\sim \\frac{1}{10} B_c$), through a succession of dynamical realignment events. In either case the net result is a laminar, but marginally stable flow, in which the original shear layer is greatly broadened. Thus, as we discussed fully in Paper I, such fields can have a remarkable stabilizing influence. This is despite the fact that their total energy content is a minor fraction of the total, so that they are nominally too weak to be important, according to the usual criteria. We considered cases in which the magnetic field was entirely within the flow plane and others in which the field was oblique to that plane, in order to examine the role in nonlinear flows of the component out of the plane. For the $2\\frac{1}{2}$ D flows we have studied, only the field components in the flow plane have any dynamical significance. In fully 3 D flows, however, we expect further evolution of the ``quasi-steady relaxed states'' of both very weak field (or dissipative) cases and weak field (disruptive) cases. The cat's eye vortex of very weak field cases is subject to a 3 D instability known as the elliptical instability (\\cite{pierr86}; \\cite{bayly86}) unless the flow lines around the vortex follow perfect circles. The planar shear flow of weak field cases is stable against linear perturbations but unstable to 3D finite-amplitude perturbations (\\cite{boh88}). Thus, it will be important to extend the present study to the fully 3 D regime, and we are preparing to do that." + }, + "9701/astro-ph9701050_arXiv.txt": { + "abstract": "In this course we review the theory of incompressible homogeneous turbulence at an elementary level, and discuss the similarities and differences expected in the compressible case, relevant to the interstellar medium and molecular clouds. We stress that a general definition of turbulence applicable to the compressible case should not rely on the Kolmogorov $k^{-5/3}$ spectrum nor on an energy cascade from large to small scales. Instead, we discuss the various possibilities for the energy spectrum of compressible turbulence, which numerical simulations suggest should be $\\sim k^{-2}$, and the nature of the cascades, if at all present. We then discuss issues concerning molecular clouds which are likely to be directly related to turbulence, such as cloud formation, cloud structure, and cloud support against gravity. ", + "introduction": "\\markboth{Enrique V\\'azquez-Semadeni}{Turbulence in Molecular Clouds} In recent years, a wide variety of scientific disciplines have come to the realization that nonlinear phenomena are the norm rather than the exception. Most physical, astronomical, biological, social and economic systems are strongly nonlinear, and in fact the time is ripe for a change in our categorization of such systems. While our current classifications of dynamical behavior are based on a ``linear vs.\\ nonlinear'' scheme, intrinsically attributing preponderance to linear systems, a ``complex vs.\\ non-complex'' classification would probably be more appropriate, reflecting the fact that most systems in Nature and society are complex. Turbulent flows are a prime example of complex systems, and it is well known that a complete theory even of incompressible turbulence does not exist. By this it is meant that some specific statistical properties of turbulent flows, such as the energy spectrum and the higher order correlation functions, cannot be derived directly from the equations of motion without the introduction of simplifying assumptions. Nevertheless, phenomenological theories exist for both three-dimensional (Kolmogorov 1941) and two-dimensional (Kraichnan 1967) turbulence in incompressible flows, which have been extremely influential. In fact, the Kolmogorov theory has almost become synonymous with turbulence. However, it will be argued in this course that such an identification is likely not to be adequate for compressible turbulence, and thus a more appropriate definition of turbulence is that {\\it a)} it contains an extremely large number of excited degrees of freedom (or modes); {\\it b)} the modes are able to nonlinearly exchange excitation; {\\it c)} the system is unpredictable in the sense of exhibiting {\\it sensitivity to initial conditions}, and {\\it d)} the system is {\\it mixing} (see, e.g., Scalo 1987; Lesieur 1990; Frisch 1995). Property c) is a distinctive feature of {\\it chaotic} systems. The interstellar medium (ISM) in particular is an extremely complex system, including gaseous, dust, cosmic ray and magnetic field components in a strongly turbulent state. As discussed in the Virial Theorem chapter of this book, the gaseous and magnetic components of the ISM are probably well described by the magnetohydrodynamic (MHD) equations (see, e..g., Cowling 1976; Spitzer 1978; Shore 1992; Shu 1992). In this course, we will discuss some basic theoretical aspects of compressible turbulence, remarking its differences with the incompressible case (\\S\\ \\ref{theory}), and then review some observational and theoretical work regarding several aspects of turbulence and molecular clouds, namely cloud formation, cloud structure and cloud support and star formation (\\S\\ \\ref{applications}), emphasizing on recent developments. Extensive reviews of earlier work can be found in Dickman (1985) and Scalo (1987). Finally, \\S 4 presents a summary and conclusions. ", + "conclusions": "In this course we have reviewed the basic theory of incompressible and compressible turbulence, and then attempted to describe the wide variety of interstellar and molecular cloud phenomena which are likely to be related to turbulence, in particular the processes of cloud formation, cloud structuring and cloud support. It should be emphasized that most treatments are phenomenological, since a full theory of interstellar turbulence is lacking. Such a theory should be able to predict properties like the filling factor of the dense regions, the density pdf, the density power spectrum and correlation function, and the density contrasts at each hierarchical level (rather than assuming them); the velocity spectrum and correlation as a function of the available energy sources, and the efficiency of star formation. However, since an equivalent theory is not available even in the incompressible case, it is clear that interstellar turbulence research is still in a highly incipient phase. \\vspace*{0.3cm}" + }, + "9701/astro-ph9701116_arXiv.txt": { + "abstract": "Proof of the existence of a significant population of normal disk galaxies at redshift $z>2$ would have profound implications for theories of structure formation and evolution. We present evidence based on Keck HIRES spectra that the damped Ly$\\alpha$ absorber at $z=3.15$ toward the quasar Q 2233+1310 may well be such an example. Djorgovski et al. have recently detected the Ly$\\alpha$ emission from the absorber, which we assume is at the systemic redshift of the absorbing galaxy. By examining the profiles of the metal absorption lines arising from the absorbing galaxy in relation to its systemic redshift, we find strong kinematical evidence for rotation. Therefore the absorber is likely to be a disk galaxy. The inferred circular velocity for the galaxy is $\\geq200$ km s$^{-1}$. With a separation of $\\simeq17$ kpc ($q_0=0.1$, $H_0=75$) between the galaxy and the quasar sightline, the implied dynamic mass for the galaxy is $\\geq 1.6\\times 10^{11} M_{\\odot}$. The metallicity of the galaxy is found to be [Fe/H]$=-1.4$, typical of damped Ly$\\alpha$ galaxies at such redshifts. However, in another damped Ly$\\alpha$ absorber at $z=2.81$ toward Q 0528$-$2505, no kinematical evidence for galactic rotation is evident. In the latter case, the damped Ly$\\alpha$ absorber occurs near the background quasar in redshift so its properties may be influenced by the background quasar. These represent the only two cases at present for which the technique used here may be applied. Future applications of the same technique to a large sample of damped Ly$\\alpha$ galaxies may allow us to determine if a significant population of disk galaxies already existed only a few billion years after the Big Bang. ", + "introduction": "The existence (or lack) of a significant population of normal disk galaxies at very high redshifts ($z>2$) would have profound implications for models of structure formation and evolution in the early universe. While the damped Ly$\\alpha$ (DLA) systems seen in spectra of high-redshift quasars have been suggested to represent disk galaxies in their youth (Wolfe et al 1986), direct evidence for such a hypothesis has been scarce (see the discussion by Lu et al 1996). Wolfe (1995) has suggested that the metal absorption lines associated with DLA galaxies exhibit what was termed ``edge-leading'' asymmetric profiles, as expected from rotating gaseous structures. Detailed analysis of a larger sample of DLA systems appears to confirm this result (Prochaska \\& Wolfe 1997). The DLA system at $z_{damp}=3.15$ in the spectrum of the background quasar Q 2233+1310 ($z_{em}=3.30$) has a neutral hydrogen column density $N$(H I)$\\simeq 10^{20}$ atoms cm$^{-2}$ (Lu et al 1993). Very recently, Djorgovski et al (1996) reported the detection of Ly$\\alpha$ emission at $z=3.1530$ only 2.3\" away from the quasar. They argued that the Ly$\\alpha$ emission most likely comes from the same galaxy that is also responsible for the DLA absorption seen in the quasar spectrum. The separation between the galaxy and the quasar sightline implies a radius of at least 17 kpc ($q_0=0.1$ and $H_0=75$) for the galaxy, comparable to the size of normal spiral disks. Djorgovski et al suggested that the offset of $\\sim 200$ km s$^{-1}$ between the redshift of the Ly$\\alpha$ emission and that of the damped Ly$\\alpha$ absorption may be due to galactic rotation, although no independent supporting evidence was available. The detection of the Ly$\\alpha$ emission provides a measure of the systemic redshift of the absorbing galaxy. An examination of the metal absorption line profiles in relation to the systemic redshift may then reveal important kinematic information about the absorbing gas and hence the nature of the absorbing galaxy. Motivated by this, we obtained an echelle spectrum of Q 2233+1310 in July 1996 using the 10m Keck telescope and the HIRES spectrograph. Details of the observations, data reductions, and full analysis of the DLA system will be described elsewhere (Lu, Sargent, \\& Barlow 1997, in preparation). Here we only discuss the results relevant to the galactic rotation hypothesis. ", + "conclusions": "By comparing the observed metal absorption line profiles with simulated model profiles, Prochaska \\& Wolfe (1997) find strong statistical evidence that DLA absorbers at $\\sim2.5$ are rotating disks with a most likely circular velocity $v_{rot}\\simeq 250$ km s$^{-1}$. The technique discussed here presents a way to test independently the hypothesis that DLA absorbers are young disk galaxies. The analysis of Prochaska \\& Wolfe (1997) has the virtue of a large sample size, but the results are only statistical. The technique presented here allows for the test of the galactic rotation hypothesis for individual cases, but is only applicable when the systemic redshift of the absorber is known. The two techniques should complement each other well in revealing the nature of DLA absorbers at $z>2$, when it becomes extremely difficult to determine the morphology of the absorbing galaxies using standard imaging techniques. The combination of rotation velocities as inferred from absorption line analysis with information on the impact parameters found from imaging studies allows for the estimate of the typical mass of the absorbing galaxies. Such results should provide significant constraints on theories of structure formation in the early universe." + }, + "9701/hep-ph9701423_arXiv.txt": { + "abstract": "We show that gravitational radiation is produced quite efficiently in interactions of classical waves created by resonant decay of a coherently oscillating field. For simple models of chaotic inflation in which the inflaton interacts with another scalar field, we find that today's ratio of energy density in gravitational waves per octave to the critical density of the universe can be as large as $ 10^{-12}$ at the maximal wavelength of order $10^{5}$ cm. In the pure $\\lambda\\phi^4$ model, the maximal today's wavelength of gravitational waves produced by this mechanism is of order $10^6$ cm, close to the upper bound of operational LIGO and TIGA frequencies. The energy density of waves in this model, though, is likely to be well below the sensitivity of LIGO or TIGA at such frequencies. We discuss possibility that in other inflationary models interaction of classical waves can lead to an even stronger gravitational radiation background. ", + "introduction": "Recent research in inflationary cosmology has attracted attention to highly non-equilibrium states created in a decay of a coherently oscillating field after the end of inflation. These states could support a number of non-equilibrium phenomena, such as non-thermal symmetry restoration \\cite{effects} and baryogenesis \\cite{effects,bau} shortly after or during the decay of the oscillating field. In this paper we want to show that non-equilibrium states produced by the decay of coherent oscillations of a field are a quite efficient source of a stochastic background of gravitational waves. There are several possible processes in the early universe capable to produce a stochastic background of relic gravitational waves. One is the parametric amplification of vacuum graviton fluctuations during inflation \\cite{LG75}. This process is efficient on all frequency scales. Waves with lowest frequencies cause inhomogeneities of cosmic microwave background \\cite{RSV}. In conventional scenarios, this restricts the amplitude of high-frequency gravitational waves to be far below \\cite{mt} the experimental limits accessible for direct detection experiments in near future ; this conclusion changes in superstring motivated cosmologies \\cite{ss}. Another source of gravitational radiation is classical emission that accompanies collisions of massive bodies. A natural source in this class in the early universe is a strongly first-order phase transition when gravitational waves are produced in collisions of bubbles of a new phase \\cite{ptr1,TW,ptr2}, in particular the phase transition that terminates first-order inflation \\cite{TW}. Gravitational radiation is also emitted during the decay of a cosmic string network \\cite{cs}. In this paper we discuss a new source of relic gravitational waves. Decay of coherent oscillations of a scalar field can produce large, essentially classical fluctuations via parametric amplification (resonance) \\cite{GMM}. These classical fluctuations, which can be viewed either as classical waves traveling through the universe or, at least qualitatively, as quantum ``particles'' in states with large occupation numbers, interact with the oscillating background and each other. This interaction, which we call rescattering \\cite{us}, is accompanied by gravitational radiation. That is the effect we want to estimate. Favorable conditions for an effective parametric resonance in cosmology naturally appear in inflationary models \\cite{KLS,resonance}. We consider two types of simple inflationary models here. One type of models has two scalar fields with an interaction potential of the form $g^2\\phi^2 X^2/2$, and the resonance produces mostly fluctuations of a scalar field $X$ other than the field $\\phi$ that oscillates (although subsequent rescattering processes produce large fluctuations of the field $\\phi$ as well). We consider a range of moderate values of the coupling $g^2$ (see below); in this case fluctuations are not suppressed too strongly by non-linear effects (cf. Refs. \\cite{wide,resc,PR}). In simplest models of this type, $\\phi$ is the inflaton itself. For these models, we find that typically $\\sim 10^{-5}$ of the total energy of the universe go into gravitational waves, at the time of their production. The minimal today's frequency $f_{\\min}$ of these waves is typically of order $10^{5}$ Hz, and today's spectral density at this frequency can be as large as $10^{-12}$ of the critical density. Another type of model we considered was the pure $\\lambda\\phi^{4}$ model of chaotic inflation. In this model, the minimal today's frequency $f_{\\min}$ is of order $10^{4}$ Hz, close to the upper bound of the operational LIGO and TIGA frequencies \\cite{LIGO}: 10 Hz $\\alt f_{\\rm LIGO}\\alt 10^4$ Hz. We do not yet have efficient means of extrapolating our numerical results for today's spectral intensity to the minimal frequency of this model, but we do not expect it to be above $10^{-11}$ of the critical density. That would be well below the sensitivity of LIGO or TIGA at frequencies of order $10^{4}$ Hz. We believe, however, that in the absence of a commonly accepted specific inflationary model, it is premature to rule out possibility of experimental detection, even already by LIGO or TIGA, of gravity waves produced by the mechanism we consider here. At the end of this paper, we discuss possibility of a stronger background of gravitational waves in models with more fields or more complicated potentials. ", + "conclusions": "We have shown that gravitational radiation is produced quite efficiently in interactions of classical waves created by resonant decay of a coherently oscillating field. For simple models of chaotic inflation in which the inflaton interacts with another scalar field, we find that today's ratio of energy density in gravitational waves per octave to the critical density of the universe can be as large as $ 10^{-12}$ at the maximal wavelength of order $10^{5}$ cm. In the pure $\\lambda\\phi^4$ model, the maximal today's wavelength of gravitational waves produced by this mechanism is of order $10^6$ cm, close to the upper bound of operational LIGO and TIGA frequencies. The energy density of waves in this model likely to be well below the sensitivity of LIGO or TIGA at such frequencies. In other types of inflationary models (or even with other parameters) the effect can be much stronger. We do not exclude that among these there are cases in which it can be observable already by LIGO or TIGA. The relevant situations are: \\begin{enumerate} \\item At some values of the coupling constant $g^2$ (or the resonance parameter $q$), the most resonant momenta are close to $k=0$. In the model with massless inflaton this happens, for example, for $q=100$ (and does not happen for $q=30$ or $q=105$, which we discussed so far; in these cases, the most resonant momenta are at $k\\sim 1$). The lowest frequency that we had in the box in our numerical simulations for $q=100$ was $f \\approx 8\\times 10^6$ Hz. At that frequency, $\\Omega_g$ for $q=100$ was almost two orders of magnitude larger than for $q=105$. In addition, the entire spectrum of gravitational waves appears to be shifted to the left with respect to the spectra shown in Fig. \\ref{fig:Fig3}. Cases when resonance is ``tuned'' to be close to $k=0$ deserve further study. Question remains, how large, in such cases, the intensity of gravitational waves can be at the horizon scale, $H_{ch}$. \\item It is important to consider in detail models of hybrid inflation \\cite{hyb}, where the oscillating field need not be the inflaton itself, and so the frequency of the oscillations may be unrelated to the inflaton mass. \\item In models where large fluctuations produced at preheating cause non-thermal phase transitions, as suggested in Ref. \\cite{effects}, domains or strings can form. A large amount of gravitational radiation can be produced in collisions of domain walls, in a way somewhat similar to how it happens \\cite{TW,ptr2} in models of first-order inflation, or in decays of a string network, cf. Refs. \\cite{cs}. In particular, in cases when domains are formed, intensity of gravitational radiation at the horizon scale, at the moment when the domain structure disappears, is expected to be much larger than in cases without domains. \\end{enumerate} We thank L. Kofman and A. Linde for discussions and comments. The work of S.K. was supported in part by the U.S. Department of Energy under Grant DE-FG02-91ER40681 (Task B), by the National Science Foundation under Grant PHY 95-01458, and by the Alfred P. Sloan Foundation. The work of I.T. was supported by DOE Grant DE-AC02-76ER01545 at Ohio State." + }, + "9701/astro-ph9701198_arXiv.txt": { + "abstract": "Analyses of QSO absorption lines are showing that HI content has evolved over the redshift range z=5 to z= 0. The 21cm line measurements of the z=0 HI content avoid several biases inherent in the absorption line technique, such as the influence of evolving dust content in the absorbers, and will produce a reliable measure to anchor theories of galaxy evolution. Examples of important questions to be addressed by local HI surveys are: (1) is there a significant population of gas-rich galaxies or intergalactic clouds that are missing from the census of optically selected galaxies? (2) is there an adequate reservoir of neutral gas to substantially prolong star formation at its present rate? and (3) are there massive objects of such low HI column density that they can have escaped detection in the ``unbiased'' HI surveys that have been conducted so far? ", + "introduction": "\\begin{figure} \\centerline{ \\psfig{figure=briggs1.PS,width=16cm} } \\caption{Cosmological density of neutral gas, incidence of CIV absorption, and comoving density of luminous QSOs as a function of redshift from QSO absorption-line statistics. {\\it Top panel.} Mean cosmological density of neutral gas, $\\Omega_g$, normalized to the critical density (Storrie-Lombardi et al 1996; Rao et al 1995 ($z=0$); Storrie-Lombardi \\& Wolfe priv. comm.) {\\it Bottom panel.} Number of CIV metal-line absorption systems per unit redshift, $n(z)$ (Steidel 1990); $z=0.3$ point from Bahcall et al 1993). Filled points from Steidel indicate rest frame equivalent widths $W_{rest}({\\lambda}1548) > 0.15$ \\AA; open points are for $W_{rest}({\\lambda}1548) >0.3 $ \\AA. Hatched areas indicate the range ($0 0.15$ \\AA) began a sharp increase. This is also a time when luminous QSOs were most abundant. These indicators testify that we are seeing substantial redistribution of gas, as witnessed by the formation of ionized metal-rich galaxy halos and the efficient fueling of active galactic nuclei. The surge in neutral gas content indicates that protogalactic gravitational potentials were deep enough that gas was confined to sufficiently high density that it was at least momentarily immune to ionization by star bursts and ionizing background radiation. This may be the epoch that disk galaxies formed as secondary infall of gas occured into galactic potentials formed in the first round of galactic bulge formation. The disk formation would be accompanied by halo enrichment, either by in situ star formation or by metal pollution of the extended halo region by winds from the new star forming regions of the disk. An alternative view is that this is also likely to be an epoch when small protogalactic lumps are merging vigorously, and star burst within the lumps would be effective at ejecting metals into an extended region that would at later times constitute the ``halo'' region of the merger product. Figure 1 summarizes only the neutral atomic and ionized gas components. A complete balance requires an accounting for all the universe's baryons, as gas is exchanged between neutral, ionized, and molecular phases, as well as the path of stellar evolution leading to the current state where far more baryons are contained in stars than in neutral gas. Although the present HI content is only ${\\sim}10$\\% of the mass in stars, there was a period at $z\\approx 2.5$ when the HI content apparently was of order half of the present stellar mass. Estimates of the mass content in ionized halos suggest that they probably contain about ten times the mass contained in the damped Lyman-$\\alpha$ absorbers at their peak (Petitjean et al 1993). This interpretation of the CIV data relies on theoretical modeling, with large uncertainties due to ionization level and carbon abundance. It is striking that the recent HST observations (Bahcall et al 1993) are consistent with no evolution in the absorption cross section presented by high column density CIV systems from $z\\approx 1.3$ (Steidel 1990) to $z\\approx 0.3$, implying that large quantities of ionized gas may still be present, either in the form of extended halos or in intergalactic clouds whose mass could far exceed the visible stellar mass in galaxies. The hydrogen neutral fraction of these ionized clouds would provide column densities of H$^0$ well below the regime ordinarily probed by 21cm line observations. At present, the integral molecular gas mass content of galaxiers appears to be roughly equal to the neutral atomic mass (cf. Kennicutt et al 1994). \\begin{figure} \\centerline{ \\psfig{figure=briggs2.PS,width=15cm} } \\caption{Density of neutral gas as a function of time compared with $\\Omega_{stars}(z=0)$. $\\Omega_g(z)$ from references in Figure 1; $\\Omega_{stars}(z=0)$ from Lanzetta et al 1995. Rising dashed curve is KTC model for increasing stellar mass; dot-dash is KTC model for declining cold gas content, adjusted as described in text to indicate only the atomic fraction. The cross hatched band indicates the range of models proposed by Pei and Fall for the true $\\Omega_g(z)$, corrected for selection effects caused by dusty damped Lyman-$\\alpha$ absorbers. } \\end{figure} Surveys of nearby volumes in the 21 cm line are important in anchoring the $z=0$ point of Figure 1. A complete inventory of neutral gas in the nearby Universe, of the sort that is being provided by unbiased 21cm line surveys, will have tight error bars and thus will carry large statistical weight in models that describe the evolution of $\\Omega_g$. Note that the statistical errors of the $z>1.6$ measurements of $\\Omega_g$ are so large that these high $z$ points are consistent with no evolution at all. In principle, the low $z$ regime ($0 < z < 1.6$) of the diagram can be measured using QSO absorption line methods in much the same way as the high $z$ points. However, the observations are difficult since the Lyman-$\\alpha$ line is not redshifted into the optical window until $z \\geq 1.6$, so the low $z$ absorption line work must be done from space observatories such as IUE and HST (Lanzetta et al 1995, Rao et al 1995). When the data points are plotted as a function of time, as in Figure 2, it is clear that the single low $z$ QSO absorption line point applies to a time span well longer than half the age of the Universe. Several selection effects make it difficult to obtain a reliable sampling of damped Lyman-$\\alpha$ absorbers at low $z$. Cosmological factors, as well as the apparent shrinking of damped Lyman-$\\alpha$ absorption cross section, $\\sigma(t)$, with increasing age of the Universe (Lanzetta et al 1995), act to make them very rare at recent times: $dN/dt \\propto \\sigma(t)t^{-k}$, where $k=2$ for $q_o=1/2$ and $k=3$ for $q_o=0$. There are also selection effects that are likely to influence the QSO absorption line measurements and that affect the low $z$ measurement most strongly: (1) The presence of dust in disk galaxies is expected to become increasingly important as they evolve and may act to selectively attenuate the light from QSOs behind them, causing these lines of sight to be under represented in QSO samples (Fall \\& Pei 1989, Pei \\& Fall 1995, Webster et al 1995) although this view is contested (Boyle \\& Di Matteo 1995). (2) Gravitational lensing may act to selectively amplify background objects into QSO samples, but also may bend the light path so as to dodge the high HI column densities of the disk (Bartelmann \\& Loeb 1996, Smette et al 1995a,1996). Figure 2 includes curves to indicate trends in stellar evolution relative to the decline in $\\Omega_g$ with time. Recent analyses of the depletion of the integral mass content of galaxies over cosmological times due to star formation have been presented by Lanzetta et al (1995) and Pei and Fall (1996). A related study by Kennicutt et al (1994, KTC) addresses the prolonging of the current star formation rate in $z\\approx 0$ disks due to delayed gas return as stellar populations age. An example of the KTC models is presented in Figure 2 to illustrate both the rise of stellar mass with time and decline of neutral gas content with time for a disk system without added inflow or allowing mass to escape. For this display, the KTC model has been scaled so that the final stellar mass and the final HI mass are consistent with the observations at $z\\approx 0$; the relative proportion of HI to H$_2$ has been adjusted for this display to vary linearly with time from $\\rho_{HI}/(\\rho_{HI}+\\rho_{H_2}) = 1$ at high $z$ to 2/3 at the present time. The slope of the KTC model at the time marked ``Now'' in Figure 2 can be compared with the steeper slope drawn to indicate the rate at which the current star formation rate would consume the present atomic hydrogen content, exhausting the supply in only $\\sim$1.5 Gyr if there were no additional reservoirs of molecular gas or contribution from delayed stellar return. The models of Lanzetta et al require a balance between the decline of $\\Omega_g$ and a rise in stellar mass; along with assumptions that stellar return occurs instantaneously as each generation of stars is formed, this constrains the mean star formation rate history of the Universe. A consequence is that their models result in an uncomfortably large number of stars at $z\\approx 0$ with low metal abundances. Pei and Fall suggest that this problem can be solved with a family of dusty models, which imply that the damped Lyman-$\\alpha$ statistics drastically underestimate the integral gas content at all redshifts. The current generation of radio telescopes are not sensitive enough to resolve this controversy by simply looking back to measure the integral HI content at $z\\approx$ 1/2. On the other hand, choosing complete samples of radio selected high $z$ quasars for background probes would remove possible selection effects by dust. At $z=0$, recent 21 cm line measurements are indicating that the bulk of the atomic hydrogen content of the nearby universe is bound into galaxies with optical counterparts (Zwaan et al 1996, Schneider 1996, Szomoru et al 1994, Henning 1995, Briggs 1990). Furthermore the normalization of seems to be well understood (cf. Zwaan et al 1996, Rao \\& Briggs 1993), although there is still concern over the normalization of even the optically determined luminosity function (cf Ellis et al 1996, Glazebrook et al 1995). Clearly the determination of the integral HI content is a measurement that the Parkes Multibeam Survey will clarify since it will be complete, unbiased by extinction and optical surface brightness, and will have well understood sensitivity limits. ", + "conclusions": "" + }, + "9701/astro-ph9701151_arXiv.txt": { + "abstract": "This poster paper illustrates the color--magnitude diagrams discussed by Piotto \\etal\\ in the preceding paper. We present CMDs for 13 clusters; and we emphasize the discovery of additional blue horizontal branches in two metal-rich clusters, and the four-mode HB of NGC 2808. ", + "introduction": "The factors that determine the morphology of the horizontal branch in globular clusters are still not well understood. Metal abundance plays an important role, with the most metal-rich clusters usually having the reddest HBs. However, a number of ``second parameters,'' such as age, have been proposed as explanations of the clusters that deviate from this rule. Here we present preliminary results from a {\\it Hubble Space Telescope} program that aims to explore connections between stellar evolution and cluster dynamics, and to investigate the CMD morphology of some clusters that are difficult to observe from the ground. The central regions of ten clusters were observed with the WFPC2 on {\\sl HST}. Some of the clusters chosen had a high central density and/or concentration; others had ultraviolet flux of unknown origin detected by {\\sl IUE}. We also include here three additional CMDs that we have derived from archival data originally taken by Yanny in another program. We present two results:\\ the discovery of additional blue horizontal branches in two metal-rich globular clusters, and the intriguing ``clumpy'' nature of the blue horizontal-branch tail of NGC 2808. ", + "conclusions": "" + }, + "9701/astro-ph9701017_arXiv.txt": { + "abstract": "In 1992 the Far-Ultraviolet Space Telescope (FAUST) provided measurements of the ultraviolet (140-180nm) diffuse sky background at high, medium, and low Galactic latitudes. A significant fraction of the detected radiation was found to be of Galactic origin, resulting from scattering by dust in the diffuse interstellar medium. To simulate the radiative transfer in the Galaxy, we employed a Monte Carlo model which utilized a realistic, non-isotropic radiation field based on the measured fluxes (at 156nm) and positions of 58,000 TD-1 stars, and a cloud structure for the interstellar medium. The comparison of the model predictions with the observations led to a separation of the Galactic scattered radiation from an approximately constant background, attributed to airglow and extragalactic radiation, and to a well constrained determination of the dust scattering properties. The derived dust albedo $a~=~0.45 ~\\pm~0.05$ is substantially lower than albedos derived for dust in dense reflection nebulae and star-forming regions, while the phase function asymmetry $g~=~0.68~\\pm~0.10$ is indicative of a strongly forward directed phase function. We show the highly non-isotropic phase function to be responsible, in conjunction with the non-isotropic UV radiation field, for the wide range of observed correlations between the diffusely scattered Galactic radiation and the column densities of neutral atomic hydrogen. The low dust albedo is attributed to a size distribution of grains in the diffuse medium with average sizes smaller than those in dense reflection nebulae. ", + "introduction": "When illuminated by the galactic interstellar radiation field, dust in the interstellar medium gives rise to scattered radiation known as diffuse galactic light (DGL). Efforts to observe the DGL at different wavelengths and attempts to derive scattering properties of the dust grains from such data have extended over the past half-century, and corresponding work directed at the far-ultraviolet spectral region has been ongoing during the past quarter-century. A detailed review of DGL studies was given by Witt (\\cite{Witt90}), while Bowyer (\\cite{B91}) and Henry (\\cite{Hen91}) reviewed in-depth the particular complications involved in the observation and interpretation of the far-ultraviolet background radiation due to several sources, including the DGL. Major challenges facing DGL studies in the far-ultraviolet include the need for observations with high diffuse-source sensitivity which provide extensive sky coverage while simultaneously offering the means to separate the flux from discrete sources, such as stars and galaxies, from that of the diffuse background. For the interpretation of DGL data, there has been a demand (\\cite{Mat93}) for multiple-scattering models which take into account the inhomogeneity of the scattering interstellar medium and the anisotropic distribution of the interstellar radiation field, which is extreme for the far-ultraviolet sky (\\cite{MH95}). The successful flight of the Far-Ultraviolet Space Telescope (FAUST) (\\cite{BSLW93}), conducted as part of NASA's 1992 ATLAS-1 shuttle mission (\\cite{CT88}), has been a response to the first of these two challenges. The photon-counting imaging detector of FAUST (\\cite{Lamp90}) allowed a separation of the diffuse background from that due to stars in the field to a level where unresolved stars contribute less than one percent of the total diffuse intensity in the field (\\cite{CSB94}). At the same time, this instrument provided sensitive measurements of the diffuse background in fields covering over 1000 square degrees in representative directions ranging from high to intermediate latitudes (\\cite{SLBW95}). Sasseen et al. (1995) showed that the dominant component in the measured far-ultraviolet background is due to starlight scattered by galactic dust by demonstrating a clear relationship between spatial power spectra of the IRAS 100$\\mu$m cirrus and the FAUST far-ultraviolet diffuse background images. An initial attempt to model the FAUST ultraviolet background and to constrain the scattering properties of the diffusely distributed dust was undertaken by Sasseen \\& Deharveng (\\cite{SD96}). This attempt, which used the simple model of Jura (\\cite{Jura79}) for high-latitude dust illuminated by a constant plane source, met with only limited success however. The model was unable to correctly reflect the anisotropy and the longitudinal variation of the intensity of the illuminating radiation field. In this paper, we are using a substantially more sophisticated model to treat the radiative transfer for the diffuse background regions observed by FAUST. This model was first introduced and described briefly by Witt \\& Petersohn (\\cite{WP94}, hereafter WP) in the context of an analysis of ultraviolet background observations made by the Dynamics Explorer 1 spacecraft (\\cite{FCF89}). The organization of this paper is as follows. In \\S 2 we will summarize the data reduction and present an initial analysis. This will establish that the measured intensities are a combination of DGL, which is dependent on the column density of the dust in the line of sight, and of roughly constant contributions from airglow and possible extragalactic diffuse background radiation. In \\S 3 we will present the WP model in greater detail then was done in WP (\\cite{WP94}). Following this, we carry out the radiative transfer analysis of the FAUST data in \\S 4, followed by a discussion of the implications of the derived dust properties for the likely processing of interstellar dust in the diffuse medium away from the galactic plane in \\S 5. In \\S 6 we will summarize our results and formulate our conclusions. ", + "conclusions": "We have analyzed the 140-180nm UV background intensity in 14 FAUST fields at intermediate and high Galactic latitudes after contributions from discrete sources such as stars and galaxies had been removed. We conclude the following: (1) We confirm that a major component of the measured UV background must be of Galactic origin based on its strong correlation with N(HI), with cosec$|b|$, and with the IRAS 100$\\mu$m background intensity. We interpret this component as DGL, produced by scattering of stellar photons by Galactic dust. (2) The slope in the FAUST DGL-N(HI) correlation of 2.5~$\\pm$~0.5 [photon units/10$^{18}$~HI ~cm$^{-2}$] is found to be steeper than that seen in earlier investigations in other Galactic regions. We explain this discrepancy as owing to the existence of a very strong Galactic latitude gradient in the illuminating radiation field in the FAUST region. (3) We found our radiative transfer model capable of matching the detailed variation of the DGL intensity from one line-of-sight to another. The most important characteristic of the model responsible for this is the incorporation of the detailed anisotropy of the interstellar radiation field in the UV and its variation with z-distance from the Galactic plane. (4) The radiative transfer analysis of the DGL component with our model resulted in a set of well-constrained scattering parameters for dust in the diffuse ISM at intermediate and high Galactic latitudes: albedo $a~=~0.45~\\pm~0.05$ and phase function asymmetry $g~=~0.68~\\pm~0.10$. (5) The dust albedo is some 50\\% lower than that commonly derived for dense reflection nebulae and star forming regions. We interpret this difference as arising from a difference in size distributions with the grains in the diffuse ISM being smaller on average. (6) A contribution of about 700~$\\pm$~200 photon units appears to be uncorrelated with Galactic lines-of-sight. We interpret this component as due to a sum of residual airglow (530 ~$\\pm$~190 photon units) and isotropic extragalactic background radiation (160~$\\pm$~50 photon units, after correction for Galactic extinction). We thank Jens Petersohn for his invaluable assistance with the WP radiative transfer model during the early stages of this work. The referee, Dr. Ken Sembach, provided a number of thoughtful and constructive comments which helped to improve the presentation of this paper, and for which we are grateful. ANW and BCF acknowledge material support from NASA LTSA Grant NAGW-3168 to The University of Toledo. \\clearpage" + }, + "9701/astro-ph9701003_arXiv.txt": { + "abstract": "This paper is the fourth one of a series whose chief objective is studying the influence of different mass spectra on the dynamical evolution of open star clusters. Results from several $N$-body calculations with primordial binaries and mass loss due to stellar evolution are presented. The models show significant differences with those for primordial binaries but no stellar mass loss presented in de la Fuente Marcos (1996b). A differential dynamical behaviour depending on cluster richness is found compared to de la Fuente Marcos (1996a). The evolution of these realistic models is very dependent on the initial mass function. Even for rich clusters, there is a dependence on the binary mass spectrum. The velocity distribution of the escapers is examined and compared with results from previous calculations. The evolution of the primordial binary population is analyzed in detail. The cluster remnant and the final binary population are also studied. Finally, some conclusions about observational properties of Open Cluster Remnants are presented. ", + "introduction": "In the last few years, a significant number of pre-main-sequence binaries and multiple systems have been discovered in young open clusters (Brandner et al. 1996; Ghez et al. 1993, 1994; Leinert et al. 1991, 1993; Padgett et al. 1996; Prosser et al. 1994; Richichi et al. 1994; Simon et al. 1992, 1993, 1995). All the surveys carried out find a binary frequency which is greater than or equal to that of the field stars in the range of separations to which they are sensitive. The observational data suggest that the binary cluster population can be interpreted in terms of different formation mechanisms. Wide binaries (with periods greater than 100 years) could originate in capture events or by a fragmentation process during the collapse of a single rotating protostar. In fact, Hartigan et al. (1994) have shown that about one third of wide pre-main-sequence binary pairs in their sample (projected separations of 400--6000 AU) are not coeval, with the less massive star usually being younger than the more massive star, suggesting that they are formed by capture or exchange (more likely) events. On the other hand, close systems must almost certainly be primordial, although their exact origin is not yet well understood. Several mechanisms have been suggested (fragmentation during the late collapse, gravitational instabilities or even orbital decay) for explaining the formation of these close primordial binaries. Moreover, observations indicate that binary stars are seen since the earliest stages of star formation, which suggests primordial mechanisms for their origin (Harjunp\\\"a\\\"a et al. 1991; Reipurth \\& Zinnecker 1993; Mathieu 1992, 1994, 1996). In any case, observational results hint that star formation produces mostly binaries. Further information about binaries in clusters can be found in Trimble (1980), Abt (1983), Reipurth (1988), Mathieu (1989), Zinnecker (1989) and an extensive recent review in Bodenheimer et al. (1993). \\hfill\\par Binaries in star clusters are of chief importance both in observational and theoretical astrophysics. In some open clusters, a certain number of stars appear above the cluster main-sequence turn-off (blue stragglers); it is currently interpreted (Wheeler 1979) as a result of stellar coalescence or extended main sequence life-times caused by mixing within the stars (Abt 1985). Another popular explanation for the origin of blue stragglers may be due to binary mass transfer (McCrea 1964). Also, runaway OB stars are interpreted as binaries escaping from star clusters (De Cuyper 1982; Sutantyo 1982; Hills 1983; Gies \\& Bolton 1986). In addition, binaries in open clusters are used to determine the distance scale (standard candles). Binaries are also the key to explaining X-ray emissions from globular and old open star clusters. From a theoretical point of view, it is clear that primordial binaries (hereafter PBs) can dominate the early stages of the dynamical evolution of star clusters (see Heggie \\& Aarseth 1992). \\hfill\\par Although binaries are heavier than the mean mass in a cluster, their masses do not remain constant throughout their life. Mass loss from stellar evolution, stellar collisions and exchange of mass from one component to another can alter significantly the binary orbit or even disrupt it. Primordial binaries which are too close will be modified or even disrupted by mechanisms such as mass loss or mass transfer during the evolution of the primary from a main sequence star to a giant. Because of their massive nature, binaries tend to occupy the inner regions of star clusters so a change in their physical parameters (mass is the most important) may affect significantly the entire dynamical evolution of the system. Their impact on the cluster evolution depends for the most part on the coupling between the characteristic time-scale for mass segregation in the cluster and that for stellar mass loss. If mass segregation can be reached before significant mass loss has occurred, these primordial binaries act as energy sources more or less continually until they become unstable by stellar evolution or escape. In the case of important stellar mass loss before equipartition, this situation may not be attained and the new evolved binaries may have no tendency to reach the inner regions of the cluster because their masses have decreased in a significant way. However, from a theoretical point of view, a smaller effect is expected than in the case of models without PBs but stellar evolution, because the single stars are also losing mass, so binaries are in any case the most massive objects. \\hfil\\par The role of stellar evolution for the dynamics of star clusters has been of interest since the end of the seventies (Angeletti \\& Giannone 1979, 1980; Angeletti et al. 1980) and there are a number of recent papers (Stod\\'o{\\l}kiewicz 1982, 1985; Terlevich 1983, 1985, 1987; Applegate 1986; Weinberg \\& Chernoff 1988; Quinlan \\& Shapiro 1990; de la Fuente Marcos 1993, 1996a (hereafter Paper II); Aarseth 1996b, c). \\hfil\\par The effect of PBs on the dynamical evolution of clusters was studied before stellar evolution. The first paper on the subject was by Aarseth (1975) and it was followed by a number of papers (Aarseth 1980; Spitzer \\& Mathieu 1980; Giannone \\& Molteni 1985; McMillan et al. 1990, 1991a, b; Murphy et al. 1990; Gao et al. 1991; Hut et al. 1992; Heggie \\& Aarseth 1992; McMillan 1993; McMillan \\& Hut 1994; de la Fuente Marcos 1996b (hereafter Paper III)). In addition, Hills (1975) gave a semi-analytical discussion about this question, Goodman and Hut (1989) pointed out the importance of PBs for the evolution of globular clusters and Leonard and Duncan (1988, 1990) carried out $N$-body simulations with primordial binaries, although their main emphasis was on escapers. Summaries of most of these studies can be found in the introductory section of McMillan et al. (1990), Gao et al. (1991) and Heggie and Aarseth (1992). \\hfil\\par On the contrary, the study of the interplay between stellar mass loss and PBs on the evolution of star clusters has only very recently become of interest. Pols and Marinus (1994) studied the binary stellar evolution in young open clusters using Monte--Carlo simulations, although their chief interest was in pure stellar evolution, and not in the dynamical one. Direct $N$-body calculations have been performed principally by Aarseth, but only a few details have been published (Aarseth 1996b, c). He has found that because of stellar evolution the fraction of binaries increases in the central regions of rich star clusters ($N$ up to 10$^{4}$). This increase in the central binary fraction is because massive single stars evolve to low-mass stars. \\hfil\\par This paper is mainly devoted to study the interplay between the mass spectrum, the PB fraction and the mass loss due to stellar evolution on the dynamical evolution of open star clusters. Moreover, we want to compare our results for the surviving binary fraction with observational data for binaries in open clusters. We are mainly concerned with the binary type for trying to answer the question about the preferential type of surviving binaries in open clusters. It is to be expected that binaries with both components being low-mass stars (late spectral types) will be preferential survivors in rich open star clusters because the time-scale for cluster disruption is larger than their characteristic time-scale for significant mass loss due to stellar evolution. However, for poorly populated open clusters, the disruption time-scale can be significantly smaller than the stellar evolution time even for moderately massive stars, so it should be possible to find binaries with massive components. Also, we are interested in comparing the present results with those from previous papers in this series (de la Fuente Marcos 1995 (hereafter Paper I); Paper II; Paper III) concerning the role of the initial mass function (hereafter IMF) on the dynamical evolution of cluster models. \\hfil\\par We have performed five runs each for a total number of stars $N$ = 100, 250, 500, 750 with five different IMFs using direct integration methods. As in previous papers, we use the same version of the Aarseth's code NBODY5 (Aarseth 1985; Aarseth 1996a). This code has become a standard in the field of star clusters simulations. Written in FORTRAN, it consists of a fourth-order predictor-corrector integration scheme with individual time steps. It utilizes an Ahmad-Cohen (1973) neighbour scheme to facilitate calculation of the gravitational forces, and handles close encounters via two-, three-, four-, and chain regularization techniques (Kustaanheimo \\& Stiefel 1965; Aarseth \\& Zare 1974; Mikkola 1985; Mikkola \\& Aarseth 1993). \\hfil\\par All the calculations have been performed on a VAX 9000/210, running under OpenVMS operating system, at the {\\it Centro de Proceso de Datos (UCM, Moncloa, Madrid)}. This machine has one CPU and its peak performance is about 100 Mflops. ", + "conclusions": "The main conclusions from this work can be summarized as follows: 1.- The inclusion of stellar evolution in cluster models with a fraction of PBs affects the overall dynamical evolution of the cluster in a uncertain way which depends strongly on the cluster richness and the IMF. 2.- The stellar evolution in clusters with small population and power-law IMF accelerates their disruption in a very significant way. 3.- The escape velocity increases with the cluster richness but the dispersion decreases when $N$ increases. Sometimes, a star can leave the cluster with high velocity. 4.- The final cluster remnant is very rich in binaries, frequently purely primordial but its composition depends strongly on the initial cluster population because of the interplay between the time-scales for cluster disruption and stellar mass loss. Binaries in remnants of poor clusters do not have any special feature in their components but binaries in rich cluster remnants have usually almost the same mass for their components and have late spectral types. Collapsed objects are almost always absent from open cluster remnants." + }, + "9701/astro-ph9701235_arXiv.txt": { + "abstract": "We have analyzed the far-ultraviolet (FUV) spectra of six elliptical and S0 galaxies in order to characterize the stellar population that produces the ultraviolet flux in these galaxies. The spectra were obtained using the Hopkins Ultraviolet Telescope (HUT) during the Astro-2 mission aboard the Space Shuttle \\it Endeavour \\rm in March 1995, and cover the spectral range from 820 to 1840~\\AA\\ with a resolution of 3~\\AA. These data, together with the spectra of two galaxies observed with HUT on the Astro-1 mission, represent the only FUV spectra of early type galaxies that extend to the Lyman limit at 912~\\AA\\ and therefore include the ``turnover'' in the spectral energy distribution below Lyman alpha. Using an extensive new grid of LTE and non-LTE synthetic spectra which match the HUT resolution and cover the relevant parameter space of temperature and gravity, we have constructed synthetic spectral energy distributions by integrating over various predicted stellar evolutionary tracks for horizontal branch stars and their progeny. When the computed models are compared with the HUT data, we find that models with supersolar metal abundances and helium best reproduce the flux across the entire HUT wavelength range, while those with subsolar $Z$ \\& $Y$ fit less well, partly because of a significant flux deficit shortward of 970~\\AA\\ in the models. High $Z$ models are preferred because the contribution from the later, hotter, post-HB evolutionary stages makes up a higher fraction of the sub-Lyman~$\\alpha$ flux in these tracks. We find that AGB-Manqu$\\acute{\\rm e}$ evolution is required in all of the fits to the HUT spectra, suggesting that all of the galaxies have some subdwarf B star population. At any $Z$ \\& $Y$, the model spectra that best match the HUT flux are dominated by stars evolving from a narrow range of envelope mass on the blue end of the horizontal branch. The Astro-1 and Astro-2 data are also the first with the resolution and signal-to-noise needed to detect and measure absorption lines in the FUV spectra of elliptical galaxies, allowing a direct estimate of the abundances in the atmospheres of the stars that produce the UV flux. We find that most absorption features in the spectra are consistent with $Z = 0.1$~$Z_{\\sun}$, significantly lower than the abundances implied by the best-fitting spectral energy distributions. However, given the strong observational and theoretical evidence for diffusion processes in the atmospheres of evolved stars, the observed atmospheric abundances may not reflect the interior abundances in the population producing the ultraviolet flux in elliptical galaxies. ", + "introduction": "An understanding of the stellar populations in elliptical galaxies will shed light upon conflicting theories in cosmology, galactic evolution, and stellar evolution. However, 25 years after the initiation of space-based observations and the consequent extension of data into the ultraviolet (UV), there are still considerable uncertainties regarding the chemical composition in elliptical galaxies and the evolution of their stellar populations. The spectra of elliptical galaxies and spiral galaxy bulges exhibit a strong upturn shortward of 2700~\\AA, dubbed the ``UV upturn.'' Characterized by the ($m_{1550}-V$) color, the UV upturn shows strong variation (ranging from 2.05--4.50 mag) in nearby quiescent early-type galaxies (Burstein et al.\\ 1988\\markcite{B88}). The strength of the UV upturn is positively correlated with the strength of Mg$_2$ line absorption in the $V$ band, in the sense that the ($m_{1550}-V$) color is bluer at higher line strengths, opposite to the behavior of optical color indices (Burstein et al.\\ 1988\\markcite{B88}). Opposing theories have been devised to explain this correlation. In one camp, Lee (1994\\markcite{L94}) and Park \\& Lee (1997\\markcite{PL97}) suggest that the UV flux originates in the low metallicity tail of an evolved stellar population with a wide metallicity distribution. In the other camp, several groups (Bressan, Chiosi, \\& Fagotto 1994\\markcite{BCF94}; Greggio \\& Renzini 1990\\markcite{GR90}; Horch, Demarque, \\& Pinsonneault 1992\\markcite{HDP92}) propose that metal-rich horizontal branch (HB) stars and their progeny are responsible for the UV flux. These different metallicity scenarios in turn argue for different ages for the stellar populations in these galaxies. Ages exceeding those of Galactic globular clusters are required in the Park \\& Lee (1997\\markcite{PL97}) model, while ages as low as 8 Gyr are allowed in the Bressan et al.\\ (1994\\markcite{BCF94}) model. In these two scenarios, the EHB stars are drawn from either tail of the metallicity distribution. However, it is also possible, indeed perhaps more likely, that the EHB stars arise from progenitors near the peak of the metallicity distribution (cf.\\ Dorman, O'Connell, \\& Rood 1995\\markcite{DOR95}), but represent a relatively rare occurrence. Observations of a bimodal HB within the metal-rich open cluster NGC~6791 suggest that this might be the case (Liebert, Saffer, \\& Green 1994\\markcite{LSG94}). If the same mechanism is at work in both the elliptical galaxies and in NGC~6791, the correlation of ($m_{1550}-V$) with the global metallicity of the galaxy might indicate that this rare path of stellar evolution becomes less so at high metallicities. \\subsection{Horizontal Branch Morphology} How can two diametrically opposed interpretations arise from the same observational data? The answer lies in the parameters that govern HB morphology, such as metallicity. In general, for stars of solar metallicity or less ([Fe/H] $< 0$), higher metallicities yield redder HBs and lower metallicities yield bluer HBs, i.e. bluer HBs have more stars on the blue side of the RR Lyrae instability strip, and redder HBs have more stars on the red side. Metallicity is the implied ``first parameter'' of HB morphology. In his spectrophotometry of 13 Milky Way globular clusters, Gregg (1994\\markcite{G94}) gives an excellent observational example of the classical effect of metallicity on HB morphology; Lee, Demarque, \\& Zinn (1994\\markcite{LDZ94}) provide theoretical examples. However, it has been known for decades that other parameters affect HB morphology (cf.\\ Lee et al.\\ 1994\\markcite{LDZ94} and references therein). In this way, the metallicity debate in elliptical galaxies is tied to the ``second parameter'' debate that has been raging in the study of globular clusters. Helium abundance also plays an important role in the stellar structure of zero-age HB (ZAHB) stars (Dorman 1992\\markcite{D92}). These stars are characterized by a helium-rich core and a hydrogen-rich envelope. At the core/envelope boundary, the density $\\rho$ and mean molecular weight $\\mu$ change such that $\\rho/\\mu$ is continuous. As the helium abundance ($Y$) in the envelope increases, the $\\mu$ gradient across the core/envelope boundary decreases, which in turn decreases the temperature gradient across this boundary. Shallower temperature gradients produce higher shell temperatures, and the nuclear reactions in the hydrogen-burning shell are highly sensitive to the temperature. At the red end of the ZAHB, the shell luminosity in these massive envelopes makes a significant contribution to the total luminosity, and so an increase in temperature results in a significant increase in the total luminosity, but does not change the envelope structure significantly. At the blue end of the ZAHB, the shell luminosity in these less massive envelopes makes a smaller contribution to the total luminosity, and so the total luminosity does not change significantly with increasing temperature, but the envelope structure does. Thus, an increase in $Y$ on the blue end of the ZAHB significantly increases T$_{\\rm eff}$ and slightly increases the luminosity, while an increase in $Y$ on the red end of the ZAHB significantly increases the luminosity and slightly increases T$_{\\rm eff}$. Across the entire ZAHB, an increase in $Y$ increases the rate of nuclear reactions, resulting in more rapid consumption of the envelope. Lee et al.\\ (1994\\markcite{LDZ94}) argue that age is the second parameter driving HB morphology (see also Sarajedini, Lee, \\& Lee 1995\\markcite{SLL95}). They also investigate the role of He abundance, CNO/Fe abundance ratio, and core rotation, but rule out these latter candidates for the second parameter because of conflicts with observations. Lee et al.\\ (1994\\markcite{LDZ94}) find that age variations of several Gyr can produce dramatic differences in HB morphology, in the sense that young HBs (age~=~11--12 Gyr) tend to be red, while older HBs (age~=~15--16 Gyr) tend to be blue. This is because the mass of a red giant at the helium flash decreases as the age of a stellar population increases, which in turn decreases the mean mass of the HB stars (Lee et al.\\ 1994\\markcite{LDZ94} and references therein). Although this scenario is true in any given system, it must be stressed that when comparing different populations, the total mass loss on the red giant branch (RGB) may not have the same mean value and dispersion, and thus differences in HB morphology cannot necessarily be attributed to differences in age alone. When more than one parameter drives the HB mass distribution, it may be difficult to separate age from these other effects. For an HB with a single metallicity, the envelope mass ($M_{env}$) determines the color location of stars on the HB. According to theory, stars with lower envelope mass are bluer (hotter) and stars with higher envelope mass are redder (cooler). Stars evolving off the HB at core helium exhaustion will then follow one of three different evolutionary scenarios, again depending upon the envelope mass. Dorman, Rood, \\& O'Connell (1993\\markcite{DRO93}, hereafter DRO93) explain these evolutionary scenarios in detail; we will briefly summarize them here. The first scenario occurs for an HB star with $M_{env}\\gtrsim 0.07$~$M_{\\sun}$ (in the case of solar $Z$ and $Y$). Such a star evolves along the asymptotic giant branch (AGB), and after reaching the thermally pulsing (TP) stage, much of the envelope is removed and the star no longer sustains a deep convective exterior. At this point it evolves rapidly to higher temperatures at fixed luminosity, and then fades along the white dwarf (WD) cooling track. An example of a classical post-asymptotic giant branch (PAGB) evolutionary track (Vassiliadis \\& Wood 1994\\markcite{VW94}) is given in Fig.~1, along with the integrated far-ultraviolet (FUV) spectrum of a population of stars entering that track at the rate of one star per year. An AGB path that might precede the PAGB evolution is also shown (DRO93\\markcite{DRO93}), but this cool, red portion of the evolution (from HB to TP) makes a negligible contribution to the integrated FUV luminosity in the HUT wavelength range. The second scenario occurs if an HB star has enough envelope mass to reach the AGB, but not enough to maintain a convective exterior all the way to the TP stage (0.03~$M_{\\sun}$~$\\lesssim M_{env} \\lesssim$~0.07~$M_{\\sun}$ for $Y =$~$Y_{\\sun}$ and $Z=$~$Z_{\\sun}$). These stars, dubbed post-early AGB (PEAGB; Brocato et al.\\ 1990\\markcite{BT90}), evolve away from the AGB before reaching the TP stage, moving to higher temperatures at a constant luminosity which is less than that in the PAGB scenario. Because of the lower luminosity, and because models with lower envelope masses leave the AGB with more remaining hydrogen fuel, the PEAGB stars have longer lifetimes. Thus they emit more UV light over those lifetimes, as compared to the stars following the PAGB tracks. Fig.~1 shows an example of PEAGB evolution (DRO93\\markcite{DRO93}) and the corresponding FUV luminosity. This particular PEAGB track exhibits helium shell flash TP ``loops'' which take place on the AGB if the envelope is more than a few hundredths of a solar mass. Although they are dramatic in the HR diagram, these short-lived features do not significantly alter the integrated spectrum. The third scenario occurs for those HB stars with such small envelope masses ($M_{env} \\lesssim$~0.03~$M_{\\sun}$ for $Y =$~$Y_{\\sun}$ and $Z=$~$Z_{\\sun}$) that they are unable ever to develop deep outer convection zones, and thus are unable to reach the AGB at all. In the shell burning phase they are referred to as AGB-Manqu$\\acute{\\rm e}$ (AGBM) stars, and they spend $\\sim 10^7$~yr at relatively high luminosity ($\\sim 100$~$L_{\\sun}$) and high temperature (T$_{\\rm eff} > 30,000$~K), after which they evolve directly to the WD cooling track. The amount of FUV light emitted from AGBM stars is much higher than that emitted from PAGB or PEAGB stars, as shown in Fig.~1, again due to their longer evolutionary time scales. Following the nomenclature of DRO93\\markcite{DRO93}, we will use the term extreme HB (EHB) for those HB stars that evolve along PEAGB or AGBM evolutionary paths. The different evolutionary paths a star may follow after leaving the HB are affected by the metallicity and helium abundance of the star. By itself, the envelope opacity at high $Z$ has a tendency to ``damp down'' the hydrogen burning shell, so that high-metallicity stars are fainter (Dorman 1992\\markcite{D92}). It is usually supposed that enhanced metal abundance is seen in stellar populations accompanied with increases in the helium fraction, according to a simple linear formula: $\\Delta Y/\\Delta Z =$~constant. Values for this constant are, however, determined from observations close to $Z = 0$ rather than $Z = Z_{\\sun}$, since what is often sought is the primordial helium abundance for cosmological studies. There is little information about the helium abundance of metal-rich populations and the subject is quite controversial (see Dorman et al.\\ 1995\\markcite{DOR95} and references therein; for a recent study, see Minniti 1995\\markcite{M95}). For high He envelope abundance ($Y \\gtrsim 0.40$), the envelope is consumed at a higher rate, because the hydrogen shell burns at higher temperatures (Dorman 1992\\markcite{D92}; Sweigart \\& Gross 1976\\markcite{SG76}). The result is that more stars, with a greater range of envelope mass on the HB, will follow the PEAGB and AGBM paths (Fagotto et al.\\ 1994c\\markcite{FBBC94c}; DRO93\\markcite{DRO93}). It must be stressed that it is the high {\\it absolute} value of the helium abundance that causes this effect (and not the increase in $Z$). High metallicity might affect the mass loss processes on the RGB. Compared to metal-poor RGB stars, metal-rich RGB stars have lower masses and larger radii at the same luminosity, and therefore lower surface gravities, which may in turn increase the rate of mass loss on the RGB (Horch et al.\\ 1992\\markcite{HDP92} and references therein). With greater mass loss, more stars will arrive on the EHB. We note that although the increase of mass loss with increasing metallicity is certainly plausible, the effect remains unsubstantiated by theory or observations. Once we understand the different physical parameters driving HB morphology, we may see how two different interpretations can arise for the correlation between Mg$_2$ strength and UV upturn strength. Arguing for the case of high metallicity, Bressan et al.\\ (1994\\markcite{BCF94}) say that more massive galaxies have higher mean and \\it maximum \\rm metallicity, and thus --- assuming $Y$ also becomes high for high $Z$ --- more stars have a high enough metallicity to follow the evolutionary paths of the PEAGB and AGBM stars. Taking the opposing view, Lee (1994\\markcite{L94}) argues that more massive galaxies have higher mean metallicity and are \\it older \\rm than less massive galaxies, and thus whatever metal-poor fraction there is in the stellar population will have bluer HBs. Measurements of FUV line strengths have been proposed as sensitive tests of the models by both camps (see also Yi et al.\\ 1995\\markcite{Y95}). In the Park \\& Lee (1997\\markcite{PL97}) model, one might expect to observe weak \\ion{C}{4} and \\ion{Si}{4} features anticorrelated with ($m_{1550}-V$). In the metal-rich scenario, one might expect strong lines and a positive correlation with ($m_{1550}-V$). The different metallicity scenarios have important implications for the ages of elliptical galaxies and in turn the age of the Universe. Bressan et al.\\ (1994\\markcite{BDF94}) demonstrated that in the high-$Z$ -- high-$Y$ scheme, the onset of UV excess serves as an age probe, since it is the most rapidly evolving feature in the spectrum of an old elliptical galaxy. In their example, Bressan et al.\\ (1994\\markcite{BCF94}) show that the onset of UV excess occurs in the galaxy rest frame at an age of $7.6\\times10^9$~yr. In the low-metallicity scheme, Park \\& Lee (1997\\markcite{PL97}) claim that giant elliptical galaxies must be at least 3 Gyr older than our own Galaxy, implying the lower limit to the age of the Universe must be 19 Gyr. Both models assume a Gaussian mass dispersion and either a Reimers (1975\\markcite{R75}) mass loss law or fixed mass loss. The ages and metallicities at which the UV upturn begins to appear are unfortunately quite sensitive to these assumptions and to the parameters. The production of hot HB stars in the globular clusters and in the Galactic field is not easily described by a simple function, as we do not know how to characterize mass loss in stars. The UV upturn therefore has little value as an age indicator until the physics driving the phenomenon is understood. In contrast, Dorman et al.\\ (1995\\markcite{DOR95}) argue that the size of the population of EHB stars --- i.e., the fraction of all HB stars in the galaxy that are UV-bright --- is the only sure deduction that can be made from the FUV radiation. They estimate, by using broadband colors derived from simple models of all evolutionary phases, that the blue FUV sources must have EHB fractions reaching $\\sim 20\\, \\%$ of the HB population, a result only weakly dependent on the $Z-$abundance of the models. They argue that such a fraction is not a ``trace population.'' Instead, it is more plausible that the UV flux is produced by a significant minority of the dominant population, comprised of stars which have undergone higher mass loss than their red HB counterparts. The reason for the UV-Mg$_2$ correlation remains unknown, but may be explained by an increase in the likelihood of high mass loss in more metal-rich (or, at least, $\\alpha-$ enhanced) stellar populations. In this picture there is no way to relate the FUV radiation to galaxy age, until mass loss in HB stars is better understood. In addition, the UV upturn is correlated with the velocity dispersion ($\\sigma_v$), and although the correlation between ($m_{1550}-V$) and Mg$_2$ index is tighter than that between ($m_{1550}-V$) and $\\sigma_v$ (Burstein et al.\\ 1988\\markcite{B88}), it is not yet confirmed that the underlying correlation (i.e. causal relationship) is with abundance. Because of the great uncertainties in the physical mechanisms responsible for the production of EHB stars, we have chosen in this paper to look at the constraints provided by the UV spectra themselves, with only very broad constraints on the overall stellar populations in the galaxies. Our conclusions are thus independent of any particular model of elliptical galaxy evolution. The limitation of this approach is that without a model there is no straightforward way to specify the distribution of stars on the zero-age HB. This would be a serious problem if there were a wide variety of FUV spectral energy distributions for normal elliptical galaxies. However, observations indicate that E galaxy FUV spectra are remarkably uniform (Burstein et al.\\ 1988\\markcite{B88}; Brown, Ferguson, \\& Davidsen 1995\\markcite{BFD95}) and consistent with a highly bimodal HB distribution (Ferguson 1994\\markcite{F94}). Thus, for most purposes we can model the hot stellar population with only two evolutionary tracks, one for a PAGB star and one for an EHB star. We then adjust $M_{env}$ for the EHB star and the fraction of the population in each type of star to achieve the best fit. In \\S 4.4 we discuss the effect of relaxing this assumption of extreme bimodality. \\subsection{Abundance Anomalies in Evolved Stars} It is not a straightforward exercise to determine the abundances in the stellar populations producing the UV light from elliptical galaxies. This is because the UV absorption line strengths in the spectra of evolved stars do not provide a direct measure of their inherent abundances. Numerous observations of HB stars and their progeny in our own Galaxy show that these stars exhibit abundance anomalies (see Saffer \\& Liebert 1995\\markcite{SL95} and references therein). Helium is usually underabundant in sdB stars while overabundant in sdO stars. Carbon and silicon are usually underabundant in both sdB and sdO stars, while the nitrogen abundance usually appears normal in both classes (Saffer \\& Liebert 1995\\markcite{SL95}; Lamontagne, Wesemael, \\& Fontaine 1987\\markcite{LWF87}; Michaud, Vauclair, \\& Vauclair 1983\\markcite{MVV83}). The FUV spectra of evolved stars observed with HUT show absorption lines that are markedly different from solar abundances or even scaled solar abundances (Brown, Ferguson, \\& Davidsen 1996\\markcite{BFD96}). These abundance anomalies can be partly explained by models of diffusion processes in evolved stars, which demonstrate that abundances can be greatly enhanced or diminished through chemical separation in the stellar envelope and atmosphere. Individual elements can either levitate or sink in the outer layers of a star, depending upon which force has a stronger effect on each element: gravity (g) or radiative acceleration (g$_{\\rm rad}$). In theory, the tendency is for heavy elements to sink in cooler stars, and for these elements to levitate in hotter stars. If g$_{\\rm rad}$ is larger than g, an element will be radiatively supported and it accumulates in the atmosphere, increasing line saturation, which in turn will decrease g$_{\\rm rad}$, until an equilibrium can be reached between g and g$_{\\rm rad}$ (Bergeron et al.\\ 1988\\markcite{BWMF88}). If g$_{\\rm rad}$ is less than g, the element will sink into the star. However, the resulting abundance decrease may or may not lead to an increase in the local value of g$_{\\rm rad}$, depending upon whether or not the lines were originally saturated or unsaturated. Thus, an equilibrium between g and g$_{\\rm rad}$ may or may not be reached (Bergeron et al.\\ 1988\\markcite{BWMF88}). Michaud et al.\\ (1983\\markcite{MVV83}) demonstrated that if the abundances of C, N, O, Ca, and Fe are all assumed to be originally at 1/200 solar abundance, one can achieve abundance enhancements of three orders of magnitude through diffusion on the timescale of 10$^8$ yr. Unfortunately, diffusion theory has not reached the level of accuracy needed to predict reliably the effects observed in evolved stars. For example, Bergeron et al.\\ (1988\\markcite{BWMF88}) calculated three models with effective temperatures of 20,000~K, 35,000~K, and 50,000~K, and with respective gravities of log~g~=~5.0, 5.5, and 6.0, all originally at solar abundance. They found that while the nitrogen abundance was at approximately solar abundance for all models, agreeing with observations, the behavior of carbon and silicon did not agree with observed trends. In the models, carbon was underabundant and increasing with T$_{\\rm eff}$, yet observations show that the carbon abundance seems to decrease with increasing T$_{\\rm eff}$ (Bergeron et al.\\ 1988\\markcite{BWMF88} and references therein). Silicon was predicted to be between 0.1 and 1 of the solar value, but silicon deficiencies by factors of 10$^4$--10$^5$ are observed for these stars. Such discrepancies might point to other processes at work, such as a weak stellar wind (Bergeron et al.\\ 1988\\markcite{BWMF88}). The bottom line is that these diffusion processes greatly complicate an absorption line analysis of an evolved population, and the strengths of the absorption lines are unlikely to reflect the inherent abundance of the population. Since the stronger metallic lines may well be saturated in stellar envelopes with very high abundances, metals in these envelopes may not be radiatively supported, allowing them to sink into the star. At the other extreme, very low abundances can be enhanced by orders of magnitude. Thus, diffusion processes tend to enhance low intrinsic abundances and diminish high intrinsic abundances. In the elliptical galaxies, line analysis is further complicated by the fact that the spectra represent composite systems with components that cover a range in temperature and gravity, which in turn affect line strengths just as abundance does. Furthermore, in the FUV, the density of lines is so large that only the very strongest lines will be clearly distinct in these composite spectra. With all of these complications in mind, an absorption line analysis alone will not answer the metallicity debate; we must look to the shape of the UV continuum as well. The shape of the continuum reflects the composite nature of these galaxies which no single temperature stellar model may duplicate (Brown et al.\\ 1995\\markcite{BFD95}), and thus reflects the evolutionary paths of the individual stars in the composite model. In this study, we use the stellar evolutionary models of DRO93\\markcite{DRO93}, Bressan et al.\\ (1993\\markcite{BFBC93}), and Fagotto et al.\\ (1994a\\markcite{FBBC94a}, 1994b\\markcite{FBBC94b}, 1994c\\markcite{FBBC94a}) to construct synthetic spectra for model stellar populations. The individual synthetic spectra which comprise the composite models are from the grids of Brown et al.\\ (1996\\markcite{BFD96}) and Kurucz (1992\\markcite{K92}), for T$_{\\rm eff} \\geq$~10,000~K and T$_{\\rm eff} <$~10,000~K respectively. The composite models which best fit the elliptical galaxy spectra draw the bulk of their flux from stars in the range 15,000~K~$\\leq$~T$_{\\rm eff}$~$\\leq$~30,000~K. ", + "conclusions": "Each individual analytical technique we have employed provides evidence for somewhat weak conclusions about the chemical composition of the stellar populations in these galaxies. Taken in their entirety, our results favor the hypothesis that the FUV light from elliptical galaxies originates in a population with high metallicity and possibly high He abundance. These conclusions appear to depend on the ratio of post-HB to EHB flux derived from the tracks. The results for the tracks with $Z_{evol} >$~$Z_{\\sun}$ are less distinguished from each other than they are from those of lower metallicity in this respect, so that more specific conclusions cannot be drawn. We cannot, however, from the FUV spectra alone, rule out the low-metallicity scenario of Park \\& Lee (1997\\markcite{PL97}). The preceding sections have demonstrated:\\\\ \\noindent 1) The HUT Astro-1 and Astro-2 E galaxy spectra are best fit by a composite EHB + PAGB star model; the FUV spectra are dominated by EHB stars following AGBM evolution.\\\\ \\noindent 2) When the flux from the PAGB component is constrained by the fuel consumption theorem, EHB stars of high $Y_{ZAMS}$ and $Z_{evol}$ are favored to make up the flux deficit near the Lyman limit.\\\\ \\noindent 3) Absorption features are consistent with $Z_{atm} \\simeq$~0.1~$Z_{\\sun}$ and may tend to increase as the UV upturn increases, but the direct constraints on metallicity from spectral lines are weak due to the low S/N of the data and the likely redistribution of elements within the EHB star atmospheres. If EHB stars alone could account for the FUV light in elliptical galaxies, then EHB stars with high $Z_{evol}$ and high $Y_{ZAMS}$ would best fit the Astro-2 data. However, we know that PAGB stars exist in ellipticals; we can observe bright PAGB stars in elliptical galaxies, and the fuel consumption theorem tells us that the stars evolving from the EHB are only a small fraction of the entire evolved population. If all of the stars in the population evolved from the EHB, the FUV light from ellipticals would be much brighter than that observed, since EHB stars are such efficient FUV emitters. Dorman et al.\\ (1995\\markcite{DOR95}) estimate $(m_{1550}-V) \\sim 0$ for systems where all of the HB stars lose sufficient mass on the RGB to become EHB and subsequently AGBM stars (i.e. sdB and then sdO stars). Hence the theoretical bluest possible ($m_{1550}-V$) color implies $\\sim 6$ times more FUV radiation than is observed in the nucleus of bluest galaxy, NGC~1399. Since we know PAGB stars make up the bulk of the FUV population and contribute to at least a fraction of the FUV light, a more realistic approach is a two-component EHB + PAGB model. If the EHB and PAGB components are fit to the HUT data without fuel consumption constraints, then the HUT data can be reproduced equally well by populations with low $Z_{evol}$ \\& $Y_{ZAMS}$ or high $Z_{evol}$ \\& $Y_{ZAMS}$. The ambiguity occurs because a ``hot'' PAGB spectrum can offset the flux deficit at short wavelengths in the low $Z_{evol}$ and $Y_{ZAMS}$ models. However, an unconstrained EHB + PAGB model is unrealistic, because in such fits the number of PAGB stars exceeds that allowed by the fuel consumption theorem. Although the EHB and unconstrained EHB + PAGB models are instructive, a more realistic model is one where the number of PAGB stars is astrophysically plausible and thus constrained by the fuel consumption theorem. When we fit these constrained EHB + PAGB models to the HUT data, the fitting favors a population evolving with high $Z_{evol}$ and high $Y_{ZAMS}$. The EHB stars that best fit the Astro-1 and Astro-2 data in any EHB or EHB + PAGB model are those that lie within a narrow distribution of envelope mass on the low-mass (blue) end of the HB. The stars in this narrow distribution evolve along AGBM paths, but the mass distribution does not include those EHB stars with the very lowest envelope masses on the HB. In addition, since the SEDs of PEAGB stars are so dissimilar to the Astro-2 data (see Fig.~1), PEAGB stars cannot comprise a significant fraction of the stellar populations in these galaxies. This result is partly to be expected from the small mass range that produces them for a given $Y$ and $Z$. However, it is possible (and plausible) that relatively low mass PEAGB stars (with $\\log L/L_{\\sun} \\sim 3.2)$ are produced in old populations through mass loss on the AGB. Such objects are in principle capable of supplying the total UV flux for the weaker UV upturn systems (with $m_{1500}-V \\gtrsim 3.6$). Our modeling implies that all of the galaxies we have studied instead have some significant contribution from a small EHB population. From this we infer that the galaxies have a bimodal HB population in general. Nesci \\& Perola (1985\\markcite{NP95}) showed that a strongly bimodal temperature distribution on the HB was needed to explain the 2200~\\AA\\ dip seen in IUE spectra of elliptical galaxies, and Ferguson (1994\\markcite{F94}) invoked a bimodal temperature distribution to explain the SEDs of the Astro-1 observations of NGC~1399 and M~31, since a uniform distribution of mass on the HB would produce a ``flatter'' spectrum (Ferguson 1994\\markcite{F94}) than observed. The open cluster NGC~6791 might be an analog of the situation in elliptical galaxies, since it shows a high-metallicity evolved population with a strongly bimodal temperature and mass distribution on the HB (Liebert et al.\\ 1994\\markcite{LSG94}). However, the existence of this cluster might present problems to the theorists in both the high-metallicity and low-metallicity camps in the elliptical galaxy composition debate. It is also possible that this cluster {\\it may not} be providing us direct information about mass loss in the single stars of metal rich systems, since several of the observed stars are binaries (E. M. Green 1996, private communication) which may have affected the evolution of the giant before core helium ignition, by tidal interaction and by mass ejection. However, bimodality occurs in globular cluster HB morphology as well (e.g. NGC~1851; Walker 1992\\markcite{WA92}), where binarism is much less common. We currently lack the understanding to explain the existence of a bimodal mass distribution in a coeval, single-metallicity population. The mass-loss mechanisms on the RGB that determine the distribution of mass on the HB are the largest source of uncertainty in the evolution of these populations. Given a bimodal HB distribution where the efficient UV emitters are on the EHB, the ($m_{1550}-V$) color should track that fraction of stars evolving from the EHB, while the (m$_{<1000}-m_{1550}$) color should track the ``characteristic'' temperature of the FUV population. The uncertainties in mass distribution on the HB means that we will need more than seven data points to understand the behavior in a color-color diagram like that shown in Fig.~7. However, these FUV colors are consistent with the scenario proposed by Ferguson \\& Davidsen (1993\\markcite{FD93}) that UV-weak galaxies have a relatively large PAGB contribution to their spectra, while the UV-strong galaxies are dominated by EHB stars and their AGBM progeny, evolving from a range of temperatures on the blue end of the HB. In the future, with more data points in such a color-color diagram, we might be able to deduce information about the HB mass distribution. Park \\& Lee (1997\\markcite{PL97}) argue that stars in the low-metallicity tail of an extended metallicity distribution are responsible for the UV upturn seen in elliptical galaxies. Bressan et al.\\ (1994\\markcite{BCF94}) believe that a high-metallicity population produces the UV flux in these galaxies, and they explore metallicities as high as $Z$~=~0.05. If the absorption features in evolved populations reflect their intrinsic abundances, then the actual situation in elliptical galaxies lies somewhere in between these two extremes, since the absorption features in the HUT data are consistent with $Z_{atm} \\simeq$~0.1~$Z_{\\sun}$. There is a weak trend in the spectra of these galaxies, such that galaxies with stronger UV upturns show stronger UV line indices. This trend is in the opposite sense to that predicted by Park \\& Lee (1997\\markcite{PL97}), but in agreement with the predictions of Bressan et al.\\ (1994\\markcite{BCF94}) and Dorman et al.\\ (1995\\markcite{DOR95}). However, diffusion processes in the envelopes of stars in evolved populations tend to enhance low intrinsic abundances and diminish high intrinsic abundances, and thus it is unlikely that $Z_{atm}$~=~$Z_{evol}$. Given this situation, we must also compare the spectral energy distributions predicted by modeling the FUV data. This comparison supports the theory that the strength of the UV upturn in elliptical galaxies is directly tied to that fraction of a high $Z_{evol}$ and high $Y_{ZAMS}$ population evolving along AGBM tracks from the blue end of the horizontal branch. Integrations over high $Z_{evol}$ and high $Y_{ZAMS}$ evolutionary tracks best fit the HUT data, and AGBM evolution is more easily produced in a high $Z_{evol}$ and high Y$_{ZAMS}$ population, providing a natural explanation for the correlation between UV upturn strength and metallicity found by Burstein et al.\\ (1988\\markcite{B88}). Of course, the large uncertainties in our analysis mean that we cannot rule out the Park \\& Lee (1997\\markcite{PL97}) hypothesis. The $\\chi^2$ statistics in our fitting favor the high $Y$ \\& $Z_{evol}$ models, but to the eye one can produce reasonable matches to the HUT data with integrations over low $Y$ \\& $Z_{evol}$ tracks. In the Park \\& Lee work, they argue that the FUV flux is dominated by a population with [Fe/H]~$<$~--0.7, but unfortunately it is difficult to discern from their paper the exact distributions in $M_{env}$ and metallicity which are used in their models. Furthermore, although Park \\& Lee stress that their purpose is not to match the actual SEDs observed in ellipticals, it seems that their synthetic spectra are too flat in the FUV, at least in comparison to the overplotted IUE data (cf.\\ their Figs. 6 \\& 10). It would be interesting to apply their models to the HUT data, in order to determine if their low metallicity models can reproduce the actual SEDs of observed ellipticals, including the flux near the Lyman limit, since our analysis indicates that the low $Y$ \\& $Z$ tracks do not reproduce the short wavelength flux needed to reproduce the HUT SEDs. In order to answer definitively the debate and disentangle the effects of age and metallicity on the HB mass distribution, data with higher S/N for a larger sample of galaxies is required. At higher S/N, it should be obvious if diffusion processes are playing a large role in the strengths of the FUV absorption lines, since observations of evolved stars show lines that deviate strongly from solar abundances or even scaled solar abundances. However, we should not restrict our study to the strengths of metallic lines. The hydrogen lines are also very sensitive to the effective temperature and surface gravity of the different HB evolutionary phases. Inspection of Fig.~1 shows that PAGB spectra have a much shallower Lyman series than that seen in AGBM spectra. Although the S/N in the HUT data was not sufficient to use the Lyman series as a stellar population probe, accurate measurements of the flux below Lyman $\\alpha$ (and thus beyond the detection capabilities of HST) will yield important information about the relative numbers of PAGB and AGBM stars in a given population. Such observations could be made with another flight of HUT or a similar successor. Fortunately, future FUV measurements of elliptical galaxies will be accompanied by the rapid progress currently seen in the calculation of synthetic stellar spectra. The near future should bring the emergence of fully line-blanketed, non-LTE model atmospheres and synthetic spectra, and with these models, another source of uncertainty will be removed from our understanding of the FUV data. Once we know the chemical composition of the FUV sources in elliptical galaxies, the resolution of the debate will have impact beyond the study of stellar evolution. Elliptical galaxies show great promise as standard candles and tracers of cosmic evolution. The understanding of the FUV flux from galaxies will also enable us to decode the information stored in the longer UV wavelengths (Dorman \\& O'Connell 1996 \\markcite{DO96}). The different composition hypotheses each argue for different ages of the Universe and different scenarios for the construction of the galaxies. Under the Park \\& Lee (1997\\markcite{PL97}) low-metallicity scenario, the more massive galaxies formed first, in contrast to the ``bottom-up'' cosmologies wherein the giant elliptical galaxies formed from the merger of smaller galaxies. In this way, the metallicity debate has consequences that extend to important issues in cosmology and galaxy evolution." + }, + "9701/astro-ph9701145_arXiv.txt": { + "abstract": "We report the discovery of large-amplitude, periodic X-ray emission from the O7 V star $\\theta^1$~Orionis~C, the central star of the Orion Nebula. Ten {\\it ROSAT} HRI snapshots of the Trapezium cluster taken over the course of 21 days show that the count rate of $\\theta^1$~Ori~C varies from 0.26 to 0.41 counts~s$^{-1}$ with a clear 15-day period. The soft X-ray variations have the same phase and period as H$\\alpha$ and \\ion{He}{2} $\\lambda 4686$ variations reported by Stahl et al., and are in anti-phase with the \\ion{C}{4} and \\ion{Si}{4} ultraviolet absorption features. We consider five mechanisms which might explain the amplitude, phase, and periodicity of the X-ray variations: (1)~colliding-wind emission with an unseen binary companion, (2)~coronal emission from an unseen late-type pre-main--sequence star, (3)~periodic density fluctuations, (4)~absorption of magnetospheric X-rays in a corotating wind, and (5)~magnetosphere eclipses. The {\\it ROSAT} data rule out the first three scenarios, but cannot rule out either of the latter two which require the presence of an extended magnetosphere, consistent with the suggestion of Stahl et al. that $\\theta^1$~Ori~C is an oblique magnetic rotator. As such, $\\theta^1$~Ori~C may be the best example of a high-mass analog to the chemically peculiar, magnetic Bp stars. ", + "introduction": "$\\theta^1$~Orionis~C (HD~37022=HR~1895) is the central star of the Trapezium cluster and the principal source of ultraviolet photons illuminating the Orion Nebula (M~42). $\\theta^1$~Ori~C also illuminates the population of proplyds in its vicinity (O'Dell \\& Wen 1994), providing some of the most direct evidence yet for circumstellar disks around low-mass pre-main--sequence (PMS) stars. It is classified as O7~V (Conti \\& Leep 1974) but is spectroscopically variable (Conti 1972; Walborn 1981). Conti \\& Alschuler (1971) and Conti (1972) also found variable inverse P Cygni emission profiles in the \\ion{He}{2}~$\\lambda 4686$ line. Recently, Stahl et al. (1993) discovered that the long-known H$\\alpha$ and \\ion{He}{2}~$\\lambda 4686$ emission variations on $\\theta^1$~Ori~C exhibit a strict, stable 15.4-day periodicity, presumably the rotation period of the star. They suggested that $\\theta^1$~Ori~C may be an oblique magnetic rotator akin to magnetic chemically peculiar B stars like the He-strong star $\\sigma$~Ori~E (B2~Vpe) where the emission lines are likely to be formed in the magnetosphere above the magnetic equator (Shore \\& Brown 1990). Walborn \\& Nichols (1994) and Stahl. et al. (1996) subsequently found \\ion{C}{4} $\\lambda\\lambda 1548, 1551$ absorption variations at high velocities that are consistent with the H$\\alpha$ period and concluded that $\\theta^1$~Ori~C possesses a corotating wind. The \\ion{C}{4} absorption strength is minimum when the H$\\alpha$ and \\ion{He}{2}~$\\lambda 4686$ emission is maximum, i.e., when the wind in our line of sight is weakest. From the phase difference between the optical emission and the UV absorption-line strength, Stahl et al. (1996) infer a rotational inclination $i \\sim 45^{\\circ}$ and a magnetic obliquity $\\beta \\sim 45^{\\circ}$. X-ray variability of $\\theta^1$~Ori~C was first reported by Ku, Righini-Cohen, \\& Simon (1982) based on {\\it Einstein} HRI observations of the Trapezium. Three deep {\\it ROSAT} HRI exposures obtained in 1991 October, 1992 March, and 1992 September show that $\\theta^1$~Ori~C varied from 0.29 to 0.45 counts~s$^{-1}$. When folded with the ephemeris of Stahl et al. (1993), the high and low X-ray count rates corresponded to phases of maximum and minimum H$\\alpha$ emission, respectively (Caillault, Gagn\\'e, \\& Stauffer 1994). ", + "conclusions": "We discuss the X-ray variability in the context of five different models: (1)~colliding-wind emission with an unseen lower-mass companion, (2)~coronal emission from an unseen lower-mass companion, (3)~periodic density variations, (4)~absorption of magnetospheric X-rays in a corotating wind, and (5)~magnetosphere eclipses. \\subsection{Colliding Winds} For the first two scenarios, we assume that $\\theta^1$~Ori~C has an unseen PMS companion. While there is no compelling evidence that $\\theta^1$~Ori~C is a binary, many high-mass stars in the Trapezium are spectroscopic binaries (Abt, Wang, \\& Cardona 1991). Stahl et al. (1996) did find that a number of $\\theta^1$~Ori~C's photospheric lines show irregular $\\sim 2$~km~s$^{-1}$ radial-velocity variations, suggesting an upper limit to the companion's mass $M \\sin i \\sim 0.27 M_{\\odot}$, assuming an O star mass $M \\sim 36 M_{\\odot}$ (Howarth \\& Prinja 1989). If 15.4~d is the orbital period, then a lower limit for the binary separation is $a \\sim 85 R_{\\odot}$. In colliding-wind models, the X-rays arise in a shock region where the wind of a hot, massive star collides with either the wind or the outer atmosphere of the companion. We calculate the expected X-ray emission from $\\theta^1$~Ori~C via this mechanism by comparing it to the well-studied example of $\\gamma$~Velorum, a Wolf-Rayet binary (O9~I + WC8) in which the WR secondary possesses a very massive wind with $\\dot{M} \\sim 8.8 \\times 10^{-5} M_{\\odot}$~yr$^{-1}$ and $v_{\\rm \\infty} \\sim 1520$~km~s$^{-1}$. Willis, Schild, \\& Stevens (1995) predict $L_{\\rm X} \\sim 10^{32}$~ergs~s$^{-1}$ resulting from the $\\gamma$~Vel colliding-wind shock region. In the case of $\\theta^1$~Ori~C, the wind is variable and published values of the terminal wind speed are discordant. Prinja, Barlow, \\& Howarth (1990) used the narrow absorption component in the \\ion{C}{4} profile and estimated $v_{\\infty} \\sim 510$~km~s$^{-1}$. Walborn \\& Nichols (1994) identified a broad, variable, high-velocity component in the \\ion{C}{4} profile and determined $v_{\\rm \\infty} \\sim 3600$~km~s$^{-1}$. In order to maximize the effect of colliding winds, we assume $\\dot{M} \\sim 4 \\times 10^{-7} M_{\\odot}$~yr$^{-1}$ and $v_{\\rm \\infty} \\sim 510$~km~s$^{-1}$ for $\\theta^1$~Ori~C; in the most optimistic case, the T Tauri companion might have $\\dot{M} \\sim 4 \\times 10^{-7} M_{\\odot}$~yr$^{-1}$ (Levreault 1988) and $v_{\\rm \\infty} \\sim 200$~km~s$^{-1}$ (Natta, Giovanardi, \\& Palla 1988). Using these wind parameters in eq.~(10) of Stevens, Blondin, \\& Pollock (1992), we find that the maximum expected X-ray luminosity of $\\theta^1$~Ori~C from colliding winds would be comparable to that observed on $\\gamma$~Vel. Although the wind parameters have been chosen to maximize colliding-wind emission, the predicted $L_{\\rm X}$ falls short of the observed X-ray variability amplitude by a factor of three. Moreover, the collision of such slow winds may not produce sufficient shocks to heat the interaction region to $T > 1$~MK. If instead we use $v_{\\rm \\infty} \\sim 3600$~km~s$^{-1}$ for the terminal wind speed of $\\theta^1$~Ori~C, the predicted X-ray luminosity is a factor of at least 25 lower because of the lower wind density in the interaction region. In this case, the sharp peak at phase 0.0 predicted by the model of Willis et al. (1995) is inconsistent with the smoothly varying sinusoidal X-ray emission observed on $\\theta^1$~Ori~C. Also, colliding winds cannot account for the smooth variation with phase of the equivalent width of the \\ion{C}{4} absorption line observed by Walborn \\& Nichols (1994) and Stahl et al. (1996). It should be noted that the Stevens et al. (1992) model assumes spherically symmetric mass loss while mass loss from T Tauri stars is often collimated in bipolar jets (e.g., Shu et al. 1994). It is also unclear how long a massive accretion disk (which drives mass loss in T Tauri stars) might survive in close proximity to an O7~V star. While there are many complicating factors one could include in the colliding-wind scenario, we think it unlikely that these complications would significantly improve the agreement between the model and the observations. \\subsection{A Low-Mass Coronal Companion} Next, we consider coronal X-ray emission from an unseen, PMS companion whose coronal X-ray emission is being eclipsed by the O-star or severely attenuated by the O-star wind at the 15.422~d orbital period. Since the amplitude of the HRI variations corresponds to $\\Delta L_{\\rm X} \\sim 2.6 \\times 10^{32}$~ergs~s$^{-1}$, this must be a lower limit to the companion's X-ray luminosity. Since the most X-ray luminous low-mass star in the entire Orion region, P1817, has $L_{\\rm X} \\sim 5 \\times 10^{31}$~ergs~s$^{-1}$ (Gagn\\'e, Caillault, \\& Stauffer 1995), it is unlikely that a low-mass star can account for the observed variations. Moreover, as has been pointed out by Stahl et al. (1996), a low-mass companion cannot account for the luminosity of the H$\\alpha$ and \\ion{He}{2} $\\lambda 4686$ variations. Given the the small-amplitude radial-velocity variations, a more massive companion (e.g., an O or early-B star) could exist if the binary orbit were in the plane of the sky ($i \\approx 0$); but, in this case, we do not expect significant X-ray variability. \\subsection{Periodic Density Fluctuations} To our knowledge, $\\zeta$~Puppis (O4~I(n)f) is the only other candidate O-type magnetic rotator. Moffat \\& Michaud (1981) report periodic H$\\alpha$ variations and propose a rotational period of 5.075~d for $\\zeta$~Pup. Bergh\\\"ofer et al. (1996) report 6\\% variations in the 0.9--2.0~keV {\\it ROSAT} PSPC flux and a $2\\sigma$ peak in the X-ray power spectrum near 16.7~hr. Contemporaneous H$\\alpha$ spectra show profile variations consistent with the Moffat \\& Michaud (1981) period of 5.075~d. The residual profile variations also indicate a weak peak in the H$\\alpha$ emission power spectrum near the X-ray period of 16.7~hr. Bergh\\\"ofer et al. interpret the X-ray--H$\\alpha$ correlation as evidence for periodic density fluctuations at the base of the wind. Interestingly, the PSPC time series does not indicate any variability at the 5.1-d rotation period. Nonetheless, the periodic H$\\alpha$ and ultraviolet variations and the possible presence of correlated X-ray--H$\\alpha$ emission in $\\zeta$~Pup and $\\theta^1$~Ori~C is noteworthy. Bergh\\\"ofer et al. (1996) interpret the 16.7~hr period as pulsations or cyclically repeating azimuthal structures. They suggest that either of these will produce periodic density variations at the photosphere which propagate through the wind, producing enhanced H$\\alpha$, and X-ray emission seen from different characteristic heights in the wind. Bergh\\\"ofer et al. (1996) propose that the emissions have the same apparent phase because the phase shift between the X-ray and H$\\alpha$ emission regions is, fortuitously, $\\sim 1$. In the case of $\\theta^1$~Ori~C, the sound crossing time through the wind is short ($t \\lesssim 1$~d out to $r \\sim 10 R_{\\star}$) compared to the 15.4-d period of $\\theta^1$~Ori~C. Consequently, a density wave resulting from non-radial pulsations might lead to apparently coherent variations. However, the lowest non-radial pulsation modes for main-sequence O stars have periods in the range 0.5--2.0~d (Baade 1986), much shorter than the observed 15.4~d period. \\subsection{Absorption in a Corotating Wind} For the next two scenarios, we assume that the 15.422~d periodicity is the rotational period of a single O star. The periodic \\ion{C}{4} and \\ion{Si}{4} profile variations seen by Walborn \\& Nichols (1994) and Stahl et al. (1996) appear to require a non-spherically symmetric, corotating wind. The tremendous torque required to maintain a corotating wind also suggests an extended magnetic field geometry. As has been pointed out by Stahl et al. and Walborn \\& Nichols, the \\ion{C}{4}, H$\\alpha$, and \\ion{He}{2} variations on $\\theta^1$~Ori~C are reminiscent of the oblique magnetic rotator $\\sigma$~Ori~E (B2~Vpe). On $\\sigma$~Ori~E, mass loss occurs preferentially along open magnetic field lines over the magnetic poles, while, at lower magnetic latitudes, field lines close inside the Alfv\\'en radius, funneling wind material towards the magnetic equator. Bolton (1994) suggests that the region of closed magnetic field lines (the magnetosphere) extends out to $R \\sim 5 R_{\\star}$. If phases of maximum mass loss on $\\theta^1$~Ori~C correspond to phases of maximum \\ion{C}{4} absorption, then we expect one magnetic pole to pass our line of sight around phase 0.5. (Stahl et al. 1996). We suggest that most of $\\theta^1$~Ori~C's X-ray emission is produced in the magnetosphere of an oblique magnetic rotator. Although O-star X-rays are generally interpreted as emission from shock regions distributed throughout the wind (e.g., Owocki, Castor, \\& Rybicki 1988), X-ray emission from $\\theta^1$~Ori~C is not typical of most single, main-sequence O~stars. First, $\\theta^1$~Ori~C is the only known O star with large-amplitude, periodic X-ray variability. Second, $\\theta^1$~Ori~C's peak X-ray activity level, $L_{\\rm X}/L_{\\rm bol} \\sim 1.8 \\times 10^{-6}$, is higher (by a factor of five) than any other single O star detected in the {\\it ROSAT} all-sky survey (Bergh\\\"ofer, Schmitt, \\& Cassinelli 1996). Third, the {\\it ASCA} SIS spectrum of the Trapezium (Yamauchi et al. 1996) indicates very high temperature plasma with $T > 20$~MK. Although the SIS cannot spatially resolve $\\theta^1$~Ori~C from surrounding lower-mass stars, the HRI images and the PSPC low-resolution spectra suggest that some of the high-temperature emission must be associated with $\\theta^1$~Ori~C. Conventional O-star shock models do not predict sufficiently fast shocks to produce such hot plasma. On the other hand, magnetically confined plasma (e.g, in coronal loops on the Sun and other late-type stars) can be heated to very high temperatures. Assuming that X-rays from the magnetosphere are being absorbed in the wind, X-ray variability would arise from varying wind absorption. The amount of excess absorbing material can be inferred from the \\ion{C}{4} excess equivalent width at high velocities around phase 0.5, $W_{1548,1551} \\sim 2.2$~\\AA\\ (Walborn \\& Nichols 1994). Assuming that the absorbing material is optically thin in \\ion{C}{4}, we find a lower limit to the excess column density $N_{\\rm C IV} \\sim 3.7 \\times 10^{14}$~cm$^{-2}$. Howarth \\& Prinja (1989) determine $N_{\\rm C IV} \\sim 9.4 \\times 10^{14}$~cm$^{-2}$ from the P Cygni profile at low velocities. If the wind and the absorbing region have approximately the same abundance and ionization, then 25--30\\% of the column density measured at phase 0.5 comes from the overlying absorption region. Since the high-velocity \\ion{C}{4} absorption is not observed at phase 0.0, can this column density difference account for the observed X-ray variability amplitude? In order to estimate the X-ray spectral shape of $\\theta^1$~Ori~C, we have analyzed the {\\it ROSAT} PSPC spectrum of the Trapezium obtained in 1991 March over 4~d spanning phases 0.36--0.63, i.e., near $\\theta^1$~Ori~C X-ray minimum. We have simulated {\\it ROSAT} HRI count rates at phases 0.0 and 0.5 in XSPEC (Arnaud 1996) by fitting the PSPC spectrum obtained near phase 0.5 and varying the column density in the overlying absorption region. The fit parameters are not well constrained, but the simulations suggest that a $\\sim 25\\%$ decrease in the wind column density would result in a $\\sim 40\\%$ increase in the HRI count rate from phase 0.5 to 0.0. More detailed modeling of spatially resolved X-ray spectra is required, but our preliminary estimates suggest that absorption in the overlying, corotating wind of an oblique magnetic rotator may be a viable mechanism for the observed X-ray variability. \\subsection{Magnetosphere Eclipses} On an O-type magnetic rotator, X-ray variability might result from eclipses of the magnetosphere and/or from varying absorption in the overlying, corotating wind. If the magnetosphere on $\\theta^1$~Ori~C extends out to many stellar radii like it does on $\\sigma$~Ori~E and if $i \\sim 45^{\\circ}$, then the X-ray variations are not likely to arise from eclipses. However, if the magnetosphere were closer to the stellar surface (i.e., 1--2~$R_{\\star}$), then eclipses could produce smooth X-ray variations. This scenario can be tested with X-ray spectra obtained at opposite phases. If the X-ray spectra do not indicate any appreciable change in the absorbing column density from X-ray minimum to X-ray maximum, then the variability is most likely due to eclipses." + }, + "9701/astro-ph9701022_arXiv.txt": { + "abstract": "In this paper we study the formation, evolution and disruption of hierarchical systems in open clusters. With this purpose, $N$-body simulations of star clusters containing an initial population of binaries have been carried out using Aarseth's NBODY4 and NBODY5 codes. Stable triples may form from strong interactions of two binaries in which the widest pair is disrupted. The most frequent type of hierarchical systems found in the cluster models are triples in which the outer star is single, but in some cases the outer body is also a binary, giving a hierarchical quadruple. The formation of hierarchical systems of even higher multiplicity is also possible. Many triple systems are non-coplanar and the presence of even a very distant and small outer companion may affect the orbital parameters of the inner binary, including a possible mechanism of significant shrinkage if the binary experiences a weak tidal dissipation. The main features of these systems are analyzed in order to derive general properties which can be checked by observations. The inner binaries have periods in the range 1 -- 10$^3$ days, although rich clusters may have even smaller periods following common envelope evolution. For triple systems, the outer body usually has a mass less than 1/3 of the binary, but is sometimes a collapsed object with even smaller mass. The formation of exotic objects, such as blue stragglers and white dwarf binaries, inside hierarchical triple systems is particularly interesting. An efficient mechanism for generating such objects is the previous formation of a hierarchical system in which the inner binary may develop a very short period during a common envelope phase, which finally results in a stellar collision. ", + "introduction": "In the past few years significant evidence for a large number of binary and multiple stellar systems in galactic clusters has been obtained from observational surveys (Mathieu et al. 1990; Mermilliod et al. 1992; Ghez et al. 1993; Mason et al. 1993; Mayer et al. 1994; Mermilliod et al. 1994; Simon et al. 1995), both for zero age main sequence and pre-main-sequence stars. For multiple systems the frequency estimated for star forming regions is up to 35 \\% and for the field is up to 20 \\%. However, the fraction of these systems detected in open clusters is smaller, although the improvement in observational techniques has increased in the past few years. Currently, the majority of multiple systems discovered in open clusters are triples and quadruples. These systems are usually highly hierarchical. Triple (or even higher multiplicity) systems are found in the Pleiades (Mermilliod et al. 1992), the Hyades (Griffin \\& Gunn 1981, Griffin et al. 1985, Mason et al. 1993), Praesepe (Mermilliod et al. 1994), M67 (Mathieu et al. 1990), and NGC 1502 (Mayer et al. 1994). The majority of binary systems observed in open clusters are thought to be primordial, but there is no preferred formation mechanism for multiple systems (dynamical or primordial) at present. Most of the multiple systems studied are hierarchical because of the intrinsic stability of these systems. The origin of observed multiple systems has not been clear since the beginning of the study of these systems. Duquennoy (1988) analyzed a sample of 17 systems (14 triples and 3 quadruples) in the solar neighbourhood. He obtained a linear correlation between the logarithm of the inner and outer binary period. This was interpreted as an indication of preferential primordial origin for these systems. From a theoretical point of view, Boss (1991) has suggested that the formation of hierarchical systems occurs during the collapse of protostellar cores. Recently, Mermilliod et al. (1994) have found significant period ratios ($X = P_{out}/P_{in} \\simeq 250$) in clusters which suggest a dynamical origin. Although the question of the origin of these hierarchical systems is far from being answered, we assume here that all the hierarchical systems formed in open clusters have a purely dynamical origin. The formation and dynamical evolution of hierarchical systems in open clusters can be studied within the context of $N$-body simulations because complex interactions between stars can readily be followed in detail by numerical methods. This approach has recently been adopted (Aarseth 1996a, Kiseleva et al. 1996, Eggleton \\& Kiseleva 1996, Kiseleva 1996). In this paper the results of almost a hundred cluster models are analyzed with the purpose of studying the formation, evolution and final destinies of hierarchical systems in clusters. These models have been obtained using direct $N$-body integration by the standard workstation code NBODY5 (Aarseth 1985, 1994) and the more recent NBODY4 (Aarseth 1996b). NBODY5 consists of a fourth-order predictor corrector scheme with individual time steps. In order to account for stellar evolution and mass loss (stellar winds and supernova events), we use the fast fitting functions of Eggleton, Fitchett and Tout (1989) and Tout (1990) for population I stars. NBODY4 is based on the so-called Hermite scheme (Makino 1991) which forms the basis of a new generation of special-purpose computers (Makino, Kokubo \\& Taiji 1993). Mass loss by stellar winds is now treated according to a modified Reimers (1975) expression (Tout 1990). Chaotic tidal motion (Mardling 1995), tidal circularization (Mardling \\& Aarseth 1996), exchange of mass (Roche overflow) in binaries and magnetic braking are also included. The models with NBODY5 have been studied on a DEC 2100 4/275 AXP system. All the calculations performed with NBODY4 were made at Cambridge on the HARP-2 computer. ", + "conclusions": "This work has allowed us to discuss various aspects of hierarchical systems in open clusters. We have found that it is, indeed, possible to reproduce some observational properties (such as the linear correlation of periods) of hierarchical systems as well as to predict some characteristics of these systems for observations. Although the results described in this paper are encouraging, it is still not clear how the fraction of primordial binaries influences the formation rate of hierarchical systems and how the tidal effects, which were only discussed briefly here, affect the stability of systems with short periods. These questions will be left for future developments." + }, + "9701/astro-ph9701058_arXiv.txt": { + "abstract": "{We consider the stability of clouds surrounded by a hotter confining medium with respect to which they are in motion, against Kelvin--Helmholtz instabilities (KHI). In the presence of cooling, sound waves are damped by dissipation. Whenever cooling times are shorter than sound crossing times, as they are in the normal interstellar medium, this implies that the instability generated at the interface of the two media cannot propagate far from the interface itself. To study how this influences the overall stability, first we derive an analytic dispersion relation for cooling media, separated by a shear layer. The inclusion of dissipation does not heal the instability, but it is shown that only a small volume around the interface is affected, the perturbation decaying exponentially with distance from the surface; this is confirmed by numerical simulations. Numerical simulations of spherical clouds moving in a surroundiing intercloud medium by which are pressure confined show that these clouds develop a core/halo structure, with a turbulent halo, and a core in laminar flow nearly unscathed by the KHI. The related and previously reported ``champagne effect'', whereby clouds seem to explode from their top sides, is cured by the inclusion of radiative losses.} ", + "introduction": "Shearing flows are ubiquitous in astrophysics. A partial list includes Herbig--Haro objects (Stone, Xu \\& Mandy 1995), disk plus hot corona accretion flows around compact objects (Liang \\& Price 1977), wind outflows from galaxies (Wang 1995), jets of all sorts and scales, bipolar flows, and even winds from the progenitor of SN 1987a (Mc Cray \\& Lin 1994). Besides these, there is the immense class of two--phase media, like the ISM (Spitzer 1978), High Velocity Clouds (Ferrara \\& Field 1994, Wolfire \\etal 1995), Broad Line Regions of AGNs, protogalaxies (Silk \\& Norman 1981), Ly$\\alpha$ clouds (Sargent \\etal 1980, Giallongo 1995), which are required to give origin to just about everything astrophysical, from stars (Shu, Adams, \\& Lizano 1987), through Globular Clusters (Vietri \\& Pesce 1995), to whole galaxies (Ikeuchi \\& Norman 1991). Two--phase equilibria are common because they arise through a universal mechanism, thermal instability of radiative media (Field 1965, Balbus 1986), but, equally universally, they are subject to the shearing instability discovered by Kelvin and Helmholtz (KHI from now on), whose potential danger to ISM clouds was noticed in computer simulations at least twenty years ago (Woodward 1976) and is discussed in textbooks (Spitzer 1978). Possible stabilization of the KHI by magnetic effects was discussed at least as early as 1955 (Michael 1955, Chandrasekhar 1961) and more recently studied in detail by Miura (1984) and, particularly, by Malagoli, Bodo and Rosner (1996). There can be little doubt that this stabilization mechanism is relevant to a large class of phenomena, including jets from AGNs and Galactic sources, and Giant Molecular Clouds. Yet, the magnetic fields in several important astrophysical situations, like Ly$\\alpha$ clouds or protogalaxies, may not yet have had the time to grow to values necessary for stabilization, while in other astrophysical situations the role played by the magnetic field is not well--established. It thus seems worth its while to consider idealized situations where magnetic fields are altogether neglected. The severity of the KHI has been put into sharp focus by numerical simulations of shearing flows around unmagnetized clouds (Murray, White, Blondin and Lin 1993, MWBL from now on) which show that, in the absence of gravity, such clouds are disrupted on a time scale comparable to the flow crossing time of the cloud. These authors considered a dense cloud moving inside a potential well which also contains a tenuous, warm phase, such that the two are in pressure equilibrium. The external gravitational field makes that the cloud speed $V$ be comparable to the warm phase sound speed $c_s$, so that the relative motion is at the boundary between subsonic and supersonic, $M\\approx 1$. Then the cloud is subject to the KHI in a regime close to the incompressible one, where the growth rate is known and large. Especially destructive, in these simulations, was the development of the KHI on the largest scales. The physical explanation for this quick destruction (Doroshkevich and Zel'dovich 1981, DZ, from now on, Nittman, Falle \\& Gaskell 1982, and Nulsen 1986) is the following: when the cloud is initially placed in a wind, a stagnation point forms ahead of the cloud, whose high pressure accelerates the wind around the cloud; by Bernoulli's theorem, the pressure of the wind is least where its speed is highest, \\ie, at the top of the cloud. This pressure minimum is overcome by the inner pressure of the cloud, which is being compressed in the direction of motion by the combined effect of the stagnation point and of its own inertia, and causes an overspilling of the cloud from its top. We shall refer to this as the {\\it champagne effect}. MWBL an DZ considered the possibility that these clouds were stabilized by self-gravity, showing that this requires a minimum mass $M_{min}$ which, compared to the cloud's Jeans mass $M_J$ is \\begin{equation} \\label{zeldovich} \\frac{M_J}{M_{min}} = 0.15 \\left(\\frac{c_s}{V}\\right)^3 \\approx 0.1 \\end{equation} (MWBL), independent of the density contrast $D \\equiv \\rho_{warm}/\\rho_{cloud}$. Hence a cloud can only choose its own death: by self--gravitational collapse, if $M \\ga M_J$, or by disruption by KHI if $M \\la M_J$. We shall refer to this result as the {\\it Zel'dovich paradox}. Is there an alternative to curing the champagne effect, without incurring the Zel'dovich paradox? MWBL and DZ considered an adiabatic fluid; let us drop this assumption, and consider a fluid subject to radiative losses. It is well known that pressure waves in such a fluid are damped more often than amplified as they trudge along (Field 1965, Balbus 1986). This suggests the following remedy: that clouds are (much) larger than the distance over which pressure waves are damped. In this way, when depressurization is occurring at the top faces of the cloud, the inner part will be slow or unable to respond to the matter outflow, and the champagne effect may be slowed down, or possibly altogether removed. That cooling times are shorter than cloud crossing times may run counter to intuition (Malagoli, Bodo and Rosner 1996), but this certainly applies to the local ISM, where typical cooling times are $\\approx 10^4\\; yr$ and sound speeds $\\approx 1\\; km\\; s^{-1}$ (Spitzer 1978), resulting in damping lengths of $\\approx 3\\times 10^{16}\\; cm$, much smaller than typical cloud radii. Thus pressure waves carrying the Kelvin--Helmholtz instability inside the cloud are damped before crossing the whole cloud. We are thus trying to stabilize the KHI in a local sense. We do not expect to find a dispersion relation for the Kelvin--Helmholtz modes which shows them to be stable. Rather, we expect to find a situation where the KHI is always present, except that it is {\\it unable to propagate far from the interface}, and thus entrain in a catastrophic fate the whole cloud. The flow of the argument and the plan of the paper are as follows: in the next Section, we derive analytically the dispersion relation for the shear layer between two idealized, radiative fluids. This problem is then tackled numerically in Sec. 3, and it is shown that the KHI is not healed, but that it is confined to a narrow region around the surface of discontinuity. In Section 4, we present numerical simulations of a realistic cloud embedded in a light wind with Mach number $\\approx 1$, for both the adiabatic and the radiative cases. It is shown here that the inclusion of radiative effects cures the champagne effect. This section also contains the simulation of a cloud for a duration far exceeding the expected KHI timescale, in an effort to show that the core/halo structure that radiative clouds develop to withstand the KHI is statistically time--independent, and stable. A brief discussion in Section 5, and a short summary in Section 6 close the paper. ", + "conclusions": "Discussion, with a solution of Zel'dovich paradox} It is well--known that, when the detailed structure of the shear layer is considered, modes with wavelengths shorter than the width of the region over which there is a significant velocity gradient are stabilized against the KHI (Chandrasekhar 1961). Accordingly, Nulsen (1982) has argued that all modes with wavelength smaller than a cloud radius are also stabilized. However, MWBL's simulations, and Fig. 7 above bear witness to the strength of the remaining instability. We thus concentrate on these, longer wavelength modes, and consider Eq. \\ref{kadimensional} on the cold, dense side of the interface. Since $k c_{-,s}\\tau_- \\ll 1$ \\begin{equation} \\label{help} Re(k_a) \\approx \\frac{N}{c_{-,s}}\\left(\\frac{N\\tau_- + q}{N\\tau_- + 3(q-1)/5} \\right)^{1/2}\\;. \\end{equation} We need to treat separately the two cases $N\\tau_- \\gg 1$ and $N\\tau_- \\ll 1$. \\subsection{ Case 1: $N\\tau_- \\gg 1$} In this case, we find \\begin{equation} Re(k_a) \\approx \\frac{Re(n)}{c_{-,s}} + \\frac{2q+3}{10}\\frac{1}{c_{-,s}\\tau_-}\\;. \\end{equation} Here we can take $Re(n) \\la k V D^{1/2}$ ($D$ is the density contrast between warm and cold phases, Eq. \\ref{def2}), the value for incompressible fluids, because we know from Fig. 3 that the real instability timescale is longer for the compressible, radiative case. Then it can be shown that the second term dominates the first one if the cloud radius $\\ell$ much exceeds $\\ell_{c1}$, where \\begin{equation} \\ell_{c1} = \\frac{10 D^{1/2}}{2q+3} 2\\pi V\\tau_-\\;. \\end{equation} From the condition of pressure balance between the two phases, and that of barely sonic motion, it can be seen that the above condition is equivalent, in order of magnitude, to $k c_{-,s} \\tau_- \\ll 1$. We then obtain \\begin{equation} \\label{ka1} Re(k_a) \\approx \\frac{2q+3}{10}\\frac{1}{c_{-,s}\\tau_-}\\;. \\end{equation} From the above we find that $Re(k_a) \\gg 1$ because of our assumption $k c_{-,s}\\tau_- \\ll 1$, so that perturbations are damped in this radius range. Thus all clouds with radius in the range \\begin{equation} \\label{range1} 2\\pi c_{-,s}\\tau_- < \\ell < 2\\pi V\\tau_- \\end{equation} are stabilized against KHI. \\subsection{ Case 2: $N\\tau_- \\ll 1$} The square root in Eq. \\ref{help} can now be approximated to yield \\begin{equation} \\label{help2} Re(k_a) \\approx \\frac{5 q}{3(q-1)} \\frac{Re(n)}{c_{-,s}} + \\frac{25(2q+3)}{9(q-1)^2} \\frac{k^2 V^2\\tau_-}{2c_{-,s}}\\;. \\end{equation} Again using for $Re(n)$ the value for incompressible fluids, we see that the second term dominates the first one for $\\ell \\ll \\ell_{c2}$, where \\begin{equation} \\ell_{c2} = \\frac{2q+3}{6q(q-1) D^{1/2}} 2\\pi V\\tau_-\\;. \\end{equation} In this case, Eq. \\ref{help2} can be approximated by retaining only the second term, and it can be shown that $Re(k_a) \\gg 1$. Again, we find that all modes with radius in the range \\begin{equation} \\label{range2} 2\\pi V\\tau_- < \\ell < \\frac{2q+3}{6q(q-1) D^{1/2}} 2\\pi V\\tau_- \\end{equation} are stabilized against the KHI. Together, Eq. \\ref{range1} and \\ref{range2} imply that all modes in the range \\begin{equation} \\label{range} 2\\pi c_{-,s}\\tau_- < \\ell < \\frac{2q+3}{6q(q-1) D^{1/2}} 2\\pi V\\tau_- \\end{equation} are stabilized against the Kelvin--Helmholtz modes. It should be noticed that the above estimate is pessimistic, because, in deriving it, we have used the instability rate for incompressible fluids, rather than the slower one for realistic, compressible and radiative fluids of Fig. 3. For a numerical evaluation, we use the standard ISM parameters (Spitzer 1978): $\\tau_- \\approx 10^4 \\; yr$, $c_{-,s} \\approx 1\\; km\\; s^{-1}$, $V \\approx 10\\; km\\; s^{-1}$, to obtain that all clouds in the range $0.03 {\\rm pc}< \\ell < 33$~pc are stabilized. In particular, our simulations use the numerical values $\\tau_- = 7 \\times 10^5 \\; yr$, $c_{-,s} = 0.4\\; km\\; s^{-1}$, $V = 8\\; km\\; s^{-1}$ which implies that the cloud size we have adopted ($\\ell=40$~pc) is well into the corresponding range of stability $2 {\\rm pc}< \\ell < 3000$~pc. We can now discuss the solution of the Zel'dovich paradox. We compare $\\ell_{c2}$ with the Jeans length for a cloud $\\ell_J$: \\begin{equation} \\frac{\\ell_{c2}}{\\ell_J} \\approx 1 \\end{equation} for the standard ISM parameters. This implies that all clouds with mass $10^{-6} M_J < M < M_J$ are stabilized. This ought to be contrasted with Eq. 1. The above argument shows that clouds are stable, at least for a time comparable to the dynamical time scale $\\tau_d$. The larger question of the persistence of this flow, \\ie, how long the cloud with the core/halo structure will persist, can be dealt with here only approximately. The persistence clearly depends upon two factors: first, the net rate (gains minus losses) of mass entrainment of the cloud by the hotter medium, which we cannot determine here but which is surely smaller than $\\la \\pi \\ell^2 \\rho_h V$, where $\\rho_h$ is the density of the warm phase. Second, the kinetic energy losses due to the acceleration (in the cloud system of reference) of the warm phase matter. Both arguments lead to a deceleration timescale $\\tau_d$ of order \\begin{equation} \\tau_d = \\frac{4}{3} \\frac{R}{V} \\frac{\\rho_c}{\\rho_h} \\approx 300 \\frac{R}{V} \\;, \\end{equation} which shows the survival time to be a few hundred times the dynamical timescale, $\\approx 3\\times 10^8\\; yr$. Most likely, this time is long enough for other factors like collisions to become important. In summary, the inclusion of radiative effects has eliminated the Zel'dovich paradox, leaving a range of $\\approx 6$ orders of magnitude of mass within which clouds are stable with respect to both self--gravity and KH modes. Another caveat that is worth discussing is that it is not at all obvious that clouds in the ISM, and the confining gas should be in thermal equilibrium, but this only strengthens our arguments. It seems in fact that most clouds are slowly cooling down, with unreplenished losses. When the equation of state softens as the pressure waves trudge along, they are damped: they put more work into compressing the ISM than is returned to them because of radiative losses. So, by considering a situation of thermal equilibrium, we have put ourselves into the least favorable conditions. That some kind of stabilization is however achieved under these circumstances seems to us a sufficiently general point worth making. The last point we wish to make is that the above discussion closely parallels that made in textbooks (Landau and Lifshitz 1987) for the development of the phenomenon of separation in incompressible fluids, whereby turbulent, rotational eddies cannot penetrate the laminar flow region, with a skin depth $\\propto k^{-1}$. This is exactly similar to the discussion above, except for the different dependence of the skin depth upon wavenumber. Still, the parallel suggests where, ultimately, the shear energy will go: in turbulence of a thin layer around the surface of separation, without disturbing the laminar flow of the remaining region, a conclusion which we cannot, formally, extrapolate neither from our linear computations, nor from our coarse numerical simulations." + }, + "9701/nucl-th9701028_arXiv.txt": { + "abstract": "We search for three-alpha resonances in $^{12}$C by using the complex scaling method in a microscopic cluster model. All experimentally known low-lying natural-parity states are localized. For the first time we unambiguously show that the $0^+_2$ state in $^{12}$C, which plays an important role in stellar nucleosynthesis, is a genuine three-alpha resonance. ", + "introduction": "Carbon, which is the fundamental basis of the chemistry of terrestrial life, is produced in red giant stars by burning the helium ash of hydrogen fusion. In order to produce carbon in stellar nucleosynthesis, the $A=5$ and $A=8$ nuclear mass-stability gaps must be bridged. Salpeter and \\\"Opik pointed out \\cite{Salpeter} that the lifetime of $^8$Be is long enough, so that the $\\alpha +\\alpha \\rightleftharpoons\\; $$^8$Be reaction can produce macroscopic amounts of equilibrium $^8$Be in stars. Then, the unstable $^8$Be can capture an additional $\\alpha$ particle to produce stable $^{12}$C. However, this so-called triple-alpha reaction has very low rate because of the low density of $^8$Be. Hoyle argued \\cite{Hoyle} that in order to explain the measured abundance of carbon in the Universe, this reaction must proceed through a hypothetical resonance of $^{12}$C, thus strongly enhancing the cross section. Hoyle suggested that this resonance is a $J^\\pi=0^+$ state at $E_{\\rm r}=0.4$ MeV (throughout this paper $E_{\\rm r}$ denotes resonance energy in the center-of-mass frame relative to the three-alpha threshold, while $\\Gamma$ denotes the full width). Subsequent experiments indeed found a $0^+$ resonance in $^{12}$C in the predicted energy region. It is the second $0^+$ state ($0^+_2$) and the second excited state of $^{12}$C. Its modern parameters $E_{\\rm r}=0.3796$ MeV and $\\Gamma=8.5\\times10^{-6}$ MeV \\cite{Ajzenberg} agree well with the old theoretical prediction. We mention here that the long lifetime of the $^8$Be ground state and the existence of the $0^+_2$ resonance in $^{12}$C at the right energy region are only parts of an incredible chain of fortunate nuclear coincidences, which makes the abundant existence of carbon and oxygen possible. For an interesting account, see Ref.\\ \\cite{Rolfs}. The aim of the present work is to explore the nature of the $0^+_2$ state in $^{12}$C. The facts that this state is very close to the three-alpha threshold, and that the alpha particle is strongly bound make it probable that the wave function of $0^+_2$ has a dominant three-alpha clustering nature. The low-lying states of $^{12}$C, including $0^+_2$, have been studied in a number of macroscopic (with structureless alpha particles) \\cite{mac,Fedorov} and microscopic \\cite{mic} three-alpha models. These models reproduce the general features of the low-lying $^{12}$C spectrum. However, all these models without exception assume three-body bound state- or two-body $^8{\\rm Be}+ \\alpha$ scattering state asymptotics for the wave functions. Thus, none of them obeys the physically correct three-body boundary condition for states above the three-alpha threshold. Those models which use $^8{\\rm Be}+ \\alpha$ asymptotics with bound-state-like $^8$Be \\cite{mic} are seemingly adequate, because the small width of the $^8$Be ground state makes its wave function very similar to a bound state wave function in a large spatial region. However, one must realize that such a model predicts the states above the three-alpha threshold, e.g.\\ $0^+_2$, to be two-body, $^8{\\rm Be}+\\alpha$, resonances. This means that currently there is no unambiguous evidence that these resonances are intrinsic states of $^{12}$C. In fact, it was speculated that the $0^+_2$ state is not a three-body resonance, but an enhancement coming from the $^{12}{\\rm C}\\rightarrow\\:$$^8{\\rm Be}+\\alpha \\rightarrow \\alpha +\\alpha +\\alpha$ sequential decay \\cite{reply}. This idea was, however, criticized \\cite{comment} by arguing that all the experimental data supported the genuine $^{12}$C nature of this state. In the present paper we use a method which is able to handle the three-body dynamics of the $0^+_2$ state correctly. Thus for the first time we can unambiguously show whether this state is a genuine three-alpha resonance in $^{12}$C. We also study other low-lying natural-parity states of $^{12}$C. ", + "conclusions": "In summary, we have studied the resonances of $^{12}$C in a microscopic three-alpha model. We used the complex scaling method, which allowed us to describe the three-body Coulomb dynamics correctly for resonances. We used three different effective nucleon-nucleon interactions and their results are consistent with each other. We have localized all experimentally known low-lying natural-parity states in $^{12}$C, although a better agreement with experiment would require major improvements of our model. For the first time we were able to unambiguously show that the $0^+_2$ state of $^{12}$C, which plays an important role in the astrophysical triple-alpha process, is a genuine three-alpha resonance." + }, + "9701/astro-ph9701170_arXiv.txt": { + "abstract": "We present a hydrodynamical code for cosmological simulations which uses the Piecewise Parabolic Method (PPM) to follow the dynamics of gas component and an N--body Particle--Mesh algorithm for the evolution of collisionless component. The gravitational interaction between the two components is regulated by the Poisson equation which is solved by a standard FFT procedure. In order to simulate cosmological flows we have introduced several modifications to the original PPM scheme which we describe in detail. Various tests of the code are presented including adiabatic expansion, single and multiple pancake formation and three-dimensional cosmological simulations with initial conditions based on the cold dark matter scenario. ", + "introduction": "Most of the present cosmological models are based on the assumption that in the universe two different kinds of matter are present: the baryonic matter, which is directly observed and forms all of the bright objects, from stars to the hot gas present in X--ray clusters, and a dark, collisionless component which accounts for most of the gravitational mass in the Universe. The evolution of this system can only be described by treating both components at the same time, looking at all of their internal processes and considering their mutual interaction. In this way, making a suitable choice of the initial conditions, one can describe the evolution of cosmological structures and estimate all of the related physical observables. This can only be achieved by using numerical simulations which allow a general description of the non--linear evolution of the structures. In particular N--body techniques (Hockney \\& Eastwood 1981; Efstathiou et al. 1985; Barnes \\& Hut 1986) have proved to be particularly effective for cosmological problems in which the dynamics is controlled by gravitational forces as in the case of the dark matter and also for baryonic structures on very large scales (more than 50$h^{-1}$ Mpc). With these methods matter is described as a set of collisionless particles and its dynamics is governed by the Boltzmann equation. The N--body approach, however, is not in general suitable for describing the behaviour of the baryonic component which is also influenced by pressure forces, heating and cooling processes. The inclusion of all these phenomena requires an enormous amount of computational resources as they act on a very wide range of scales. The thermal input from gravitationally induced shocks is, for example, maximally effective at about 10 Mpc, while cooling processes are most important on scales less than 0.1 Mpc. Therefore to take into account all of these processes requires a minimum of $10^6$ resolution elements in three dimensions. The lack of adequate computational resources has delayed the development of hydrodynamic cosmological codes and only in recent years have a number of numerical schemes been proposed for following the evolution of baryonic matter. A first family of these techniques derives directly from the N--body methods. It is called ``Smoothed Particle Hydrodynamics'' (SPH) and represents fundamental fluid elements in terms of particles. The SPH methods are intrinsically Lagrangian and, following the fluid elements in their motion, have high spatial resolution and give an accurate description of high--density regions, where particles tend to concentrate. On the other hand they cannot treat properly low--density regions where few particles are present and mass resolution is poor. For a detailed introduction to these methods we refer to Hernquist \\& Katz (1989), Evrard (1990) and Steinmetz \\& Muller (1993). Recently Gnedin (1995; see also Gnedin \\& Bertschinger 1996) has presented a new approach to cosmological hydrodynamics called SLH (softened Lagrangian hydrodynamics) which utilizes a high resolution Lagrangian hydrodynamics code combined with a low resolution Eulerian solver. Its properties are intermediate between the Eulerian and SPH approaches. Multi--dimensional hydrodynamic codes not based on SPH, usually adopt an Eulerian approach and dynamical equations are solved on a fixed (or adaptive) grid. Mean values of the fluid quantities are computed in each grid cell by solving the equations of conservation of matter, momentum and energy density once the equation of state for the matter is given. Eulerian methods have good mass resolution, and can describe low--density regions better than SPH but they are spatially limited by the cell size. An Eulerian approach has been adopted in the majority of the hydrodynamical codes developed for studying large scale structures (Chiang, Ryu \\& Vishniac 1989; Cen 1992; Ryu et al. 1993; Bryan et al. 1995; Quilis, Ib\\'anez \\& Saez 1996; Sornborger et al. 1996). For a comprehensive comparison between the Lagrangian and Eulerian approaches we refer to Kang et al. (1994). In this paper we describe a cosmological hydrodynamic code that we have developed using the Piecewise Parabolic Method (PPM) introduced by Colella \\& Woodward (1984). This is a higher order extension of Godunov's shock capturing method (Godunov 1959; 1961). It is at least second--order accurate in space (up to the fourth--order, in the case of smooth flows and small timesteps) and second--order accurate in time. The high accuracy of this method allows minimization of errors due to the finite size of the cells of the grid and leads to a spatial resolution close to the nominal one. In a cosmological framework, the basic PPM technique has to be modified to include the gravitational interaction and the expansion of the universe. In the gravitational collapse of cosmological structures, extremely supersonic flows are generated and so the thermal energy cannot be computed accurately from the total energy, as is usually done. Thermal energy is only a small fraction of the total energy and so negative values or a spurious heating can both be found as a result of numerical errors. This difficulty has been overcome by computing simultaneously the total and the internal energy. The PPM algorithm has already been used for building a cosmological code by Bryan et al. (1995), however our approach differs in several respects from theirs. In the Bryan et al. work, the one--dimensional time integration is done by first performing a Lagrangian step and then remapping the results onto an Eulerian grid expressed in the usual coordinates comoving with the mean Hubble flow. We instead adopt a single--step Eulerian scheme. The construction of the effective left and right states for the Riemann problem is then more complicated than in the Lagrangian case, since the number of characteristics reaching the edge of a zone is not constant. On the other hand, this choice allows us to include in the characteristic equations both the gravitational interaction and the expansion of the universe and then the effect of all source terms is accounted for to second--order. Moreover, an Eulerian approach seems to be preferable in the case of complicated three--dimensional structures (Colella \\& Woodward 1984). The hydrodynamical part has been coupled to a Particle Mesh (PM) N--body code that describes the evolution of the dark component. The standard PM code has been modified to work with a non--constant timestep equal to that used in the hydrodynamical integration. The coupling is obtained by calculating the gravitational field due to both the components by the usual FFT procedure. The plan of the paper is as follows. In Section 2 we present the basic equations written in comoving coordinates. In Section 3 we describe the basic numerical method and the modifications needed for cosmological calculations. The numerical tests and results from cosmological simulations are presented in Section 4, while conclusions are drawn in Section 5. A more detailed description of the numerical scheme and of the method used for computing the fluxes are presented in the appendices. ", + "conclusions": "We have presented a new numerical code developed for studying the formation and evolution of cosmological structures in both baryonic and dark components. Collisional matter is treated as a fluid and the corresponding hydrodynamic equations are solved using the PPM scheme on a fixed Eulerian grid. We have described the changes to the basic method required by the cosmological applications. Particular care has been taken in including expansion and gravity in the Riemann solver and in the final integration step. This has required the calculation of the characteristic form of the hydrodynamic equations in expanding coordinates. A double formulation of the energy equation has allowed a proper treatment of the highly supersonic flows common in cosmological simulations. The behaviour of the dark matter is described using a standard Particle Mesh N--body technique, modified to allow the use of a variable timestep, as desirable for hydrodynamics. The two components are coupled through the gravitational interaction and the gravitational field is calculated from the Poisson equation using an FFT procedure. We have presented a series of tests selected for their relevance in cosmological applications, paying attention both to the accuracy of the highest resolution results and to the convergence of the method when lower resolutions are used. The one--dimensional tests show that the code can reproduce properly the expected solutions, even when very low resolution is adopted. In particular we present the results of the shock tube test and of single and multiple pancake formation. The CDM test results can be compared with those presented by Ryu et al. (1994), Kang et al. (1994) and Gnedin (1995) showing a good agreement. All of these tests have demonstrated that our code can be considered a reliable and useful tool for cosmological studies." + }, + "9701/astro-ph9701036_arXiv.txt": { + "abstract": "We present here the results of applying a new chemo-evolutionary stellar population model developed by ourselves in a previous paper (Vazdekis {\\it et al.} 1996) to new high quality observational data of the nuclear regions of two representative elliptical galaxies and the bulge of the Sombrero galaxy. Here we fit in detail $\\sim$20 absorption lines and 6 optical and near-infrared colors following two approaches: fitting a single-age single-metallicity model and fitting our full chemical evolutionary model. We find that all of the iron lines are weaker than the best fitting models predict, indicating that the iron-abundance is anomalous and deficient. We also find that the Ca$_{I}$ index at $4227 {\\rm \\AA}$ is much lower than predicted by the models. We can obtain good fits for all the other lines and observed colors with models of old and metal-rich stellar populations, and can show that the observed radial gradients are due to metallicity decreasing outward. We find that good fits are obtained both with fully evolutionary models and with single-age single-metallicity models. This is due to the fact that in the evolutionary model more than 80\\% of stars form within 1.5~Gyr after the formation of the galaxies. The fact that slightly better fits are obtained with evolutionary models indicates these galaxies contain a small spread in metallicity. ", + "introduction": "The study of the stellar populations and the distribution of metallicities plays an important role in our understanding of the star formation history of the galaxies. Their stellar populations are expected to be more complex than those of, for example, globular clusters, which are thought to be composed of a single stellar population. In fact, Burstein {\\it et al.} (1984) found differences when comparing colors and spectroscopic features of globular clusters with galaxies. Key parameters in the interpretation of the observed colors and the line-strengths are the metallicity and the age. The problem is that even in the simplest unresolved stellar systems their effects are very difficult to separate using only colors (O'Connell 1986, Renzini 1986, Buzzoni {\\it et al.} 1992). Using colors together with absorption lines more accurate conclusions can be drawn. Although every absorption line strength is dependent on different kinds of stars, in principle it should be possible to determine average metallicities or ages by carefully selecting features which more sensitive to the metallicity and others which are more sensitive to the age (e.g. Worthey {\\it et al.} 1992). However the abundances of some elements may well evolve differently from those of others (e.g. $\\alpha$-enhancement), and the conversion of ages and metallicities through models to observed colors and indices may be not unique due to problems in e.g. stellar evolution theory. Finally, the large velocity broadening in giant elliptical galaxies implies that only the strongest lines can be used to obtain physical information from their spectra. In the process of understanding the stellar population of early-type galaxies we first developed a new spectrophotometric model, which can be used to interpret observed colors and absorption lines of galaxies (Vazdekis {\\it et al.} 1996, hereafter Paper I). The model is based on the latest improvements in stellar evolution theory and on the most recent stellar libraries. Instead of studying a large sample of galaxies using a few lines indices, as has been done before (e.g. Worthey {\\it et al.} 1992, Gonzalez 1993), we preferred to obtain high quality observations of three representative early-type galaxies (two giant ellipticals and the bulge of the Sombrero galaxy), but in many colors and absorption lines, and to make very detailed fits to each index, to understand better global ages and metallicities, and also to follow the abundances of individual elements. Such analysis now is possible, since we could calibrate our observations using the large sample of stars from the extended Lick-system (Worthey {\\it et al.} (1994) hereafter WFGB). In this paper we have applied our spectrophotometric population synthesis model following both the single-age single-metallicity and the chemical evolutionary approaches. We address here the problem of whether the conclusions we obtain depend on the stellar population synthesis method we use. In the end we find that the use of many indices does yield interesting information, and we show that we can learn more than by using only a few indices, as has been done in the past. At the same time we study the stellar population gradients in the three galaxies. This paper is organized as follows: in Section~2 we explain our observations and the method we use to derive the line-strengths. In Section~3 we fit our population synthesis model and discuss the results obtained by fitting the data. Finally in Section~4 we present our conclusions. ", + "conclusions": "We have obtained high quality observations of almost the whole set of line-indices of the extended Lick-system for three representative early-type galaxies, and have applied a new spectrophotometric chemo-evolutionary population synthesis model developed by ourselves in a previous paper (Paper I). We can make models which give good fits in all the colors and many of the most important line-indices. These fits however cannot synthesize quantitatively a number of lines primarily from Fe and Ca. We find that 6 independent Fe lines are too weak compared to lines of all other elements indicating that the iron-abundance is anomalous and deficient in the radial range of the galaxies that we studied. This implies that the global metallicity inferred must depend on whether we use magnesium or iron lines as the prime indicators. Invoking $\\alpha$-enhancement one can obtain better fits for the iron lines, but other features such as NaD then become worse if we follow the abundance ratios given in Weiss {\\it et al} (1995). Finally we find that the Ca4227 is much fainter than predicted by the models. In general we find that the three galaxies require metallicities higher than solar for the inner regions while the ages are older than 10~Gyr, and the observed radial gradients are due to metallicity decreasing outward. We also find that NGC~4472 is more metal-rich, than the other two galaxies. To fit this set of galaxies with the full chemical evolutionary population synthesis model we used the variable IMF scenario (defined in Paper I) which invokes an IMF skewed towards high-mass stars in the begining, during a short period of time (smaller than 1~Gyr.), and towards low-mass stars later for the remaining time. The best fits indicate that dwarfs contribute $\\sim 70\\%$ to the U band, $\\sim 50\\%$ in V band and $\\sim 25\\%$ in K band. We find that slightly better fits are being obtained with the chemical evolutionary model than with the single-age, single-metallicity model, justifying the extra complications. However, since the predicted spread in metallicities is not very large, and since the bulk of the stars were formed at the very early stages of the galactic evolution (at age lower than $\\sim1.5~Gyr$) we conclude that the single-age single-metallicity stellar population models offer reasonable first order fits to this kind of stellar systems, especially if one wishes to avoid comprobational complexity. This study shows that it would be useful to extend the present analysis to include other features at shorter wavelengths in the UV region such as the indices of Rose (1994), and to the near-IR with indices such as Na$_{I}$ at $8190~{\\rm \\AA}$, the Ca$_{II}$ triplet and the CO or H$_{2}$O features. To understand the stellar populations of the early-type galaxies and to e.g. disentangle age and metallicity (Jones \\& Worthey 1995), Bressan {\\it et al.} 1995) it will be important to introduce as many constraints as possible, by observing the galaxies in many calibrated absorption lines. The observational data presented in this paper are available on the AAS CD-ROM, and from the WWW-homepages of the authors." + }, + "9701/astro-ph9701126_arXiv.txt": { + "abstract": "Dynamical friction and tidal disruption are effective mechanisms of evolution of globular cluster systems, especially in non--axysimmetric galaxies with a central compact nucleus. With a semi--analytical approach based on the knowledge of the dependence of the dynamical friction and tidal disruption effects on the relevant parameters, we are able to follow the time evolution of the globular cluster system in a model of a triaxial galaxy and give its observable properties to compare with observational data. An important result is that the flatter distribution of the globular cluster system relatively to that of the stellar bulge, as observed in many galaxies, can be explained by the evolution of the globular cluster system, starting from the same density profile. ", + "introduction": "Two observational facts are well established, by now: $i)$ the first is that many galaxies have globular clusters systems (GCSs) with density profiles less concentrated than their parent galaxy halo light, M87 and M49 being the prototypes (Lauer \\& Kormendy 1986, Harris 1986,1991). This has been recently confirmed by observations of a sample of 14 elliptical galaxies, made with the WFPC2 of the Hubble Space Telescope (Forbes et al. 1996). These observations, thanks to the high resolution of the HST, were able to probe the inner kiloparsecs of those galaxies and show that most, if not all, of them have GCSs with surface density profiles that rise towards the centre less steeply than the underlying galaxy light; $ii)$ the second is that many galaxies host compact massive nuclei in their centres (Dressler \\& Richstone 1988, Kormendy 1988, Kormendy \\& Richstone 1995, Eckart \\& Genzel 1996) with estimated masses in the range from 2$\\cdot 10^6 M_\\odot$ for our Galaxy and M32 up to 3$\\cdot 10^9 M_\\odot$ for M87. An explanation for the difference between halo star and globular cluster distributions has been proposed by Harris \\& Racine (1979), Harris (1986), and Racine (1991) as a difference in the formation epoch of the two components. In this scenario, the GCS formed in an earlier phase of the protogalactic collapse, while the stars that constitute the halo condensed later, this resulting in a less peaked distribution for the clusters. This picture requires an exact timing in the sequence of the evolution, in order to permit the clusters to be less metal rich than the halo stars, while producing the required differences in central concentration. In disk galaxies,however, 'the epoch of cluster formation would be early enough to force chenical enrichment but not early enough to take on a distinct spatial structure' (Harris 1986, Sect. VII, p. 840). Moreover, this scenario does not explain why the tails of the two density distributions are almost the same. This last observational evidence suggests an alternative explanation: the cluster system and the halo formed at the same time with a similar spatial distribution, and the present differences are a consequence of the dynamical evolution of the GCS. Dynamical evolution correlates also with the possible presence of masssive central nuclei. Actually, the larger core radius of the GCS would imply that the globular cluster population has been significantly depauperated in the inner regions of the system. This is the case when a massive object (like a black hole) resides in the centre of a galaxy and disrupts, by means of tidal forces, globular clusters which pass sufficiently close to it. There is no direct evidence that the aforementioned massive objects are black holes (they could be, as proposed by Kormendy \\& Richstone (1995), massive clusters of low-mass stars, stellar remnants etc.) except in the case of NGC 4258, where the discovery of a perfect keplerian rotation curve in the inner regions (Miyoshi et al. 1995) rules out, on dynamical evidences (Mayoz 1995), alternatives to black holes. In the outer regions, on the contrary, the influence of the central massive object will be negligible. So we expect that there the cluster distribution has remained more or less unchanged. Another major dynamical effect, the dynamical friction of field stars on globular clusters, enhances the efficiency of the depleting mechanism. It acts reducing the cluster orbital energy and bringing the clusters towards the centre, thus increasing the number of globular clusters in the inner regions. These clusters could feed the massive object. This scenario has been proposed first by Tremaine, Ostriker and Spitzer (1975), in a study on M31, where they showed how a massive object of mass in the range 10$^7$--10$^8$ $M_\\odot$ could directly form from globular clusters braked to the centre of the galaxy and there merged. Both of these effects (dynamical friction and tidal disruption) are significantly emphasized if the galaxy is triaxial in shape, a possibility supported by several observations (see, for example, Bertola, Vietri \\& Zeilinger (1991) which show evidence of triaxial distributions in 32 galaxies). The orbits which constitute the bulk of such a potential are the 'box' orbits, which are dense around the centre (see, e. g., Binney \\& Tremaine 1987, hereafter BT) where the massive object lies (so that even globular clusters of large apocentric distance will possibly be disrupted) and the field star density is higher (so that dynamical friction efficiency is increased). In fact, Pesce, Capuzzo Dolcetta and Vietri (1992) have demonstrated that dynamical friction decay times on box orbits are significantly reduced. The shape of the velocity ellipsoid of halo stars and globular clusters in our galaxy supports this picture. In fact, while the velocity dispersion of the halo stars is larger in the radial direction, as expected from numerical simulation of the radial collapse of the protogalaxy, the globular clusters' velocity ellipsoid is almost spherical. Under the hypothesis of a coeval formation, this can be explained by a selective process which destroyed the clusters on low-angular momentum orbits (i.e. box orbits in a triaxial galaxy). The role of triaxiality in the tidal disruption mechanism has been quantitatively discussed by Ostriker, Binney and Saha (1989) (hereafter OBS). In their paper they proposed, for the first time, that the formation of a massive nucleus from decayed globular clusters could be a self-limiting process, due to the inverse proportionality of the tidal disruption timescale, $\\tau_{tid}$, to the nucleus mass. Capuzzo Dolcetta (1993) has investigated thoroughly the evolution of a GCS in the Schwarzschild's (1979) triaxial non-rotating model, under the combined effects of dynamical friction and tidal disruption, studying the growth of the nucleus and the evolution of globular clusters' mass function. Indeed, he found that the cooperation of these two effects may lead to the formation of a compact nucleus, in form of globular clusters decayed to the centre of the galaxy. The growth of the nucleus eventually halts when its mass is large enough to shatter all incoming clusters. Of course the value of the mass reached by the growing nucleus depends, in this scheme, on the initial GCS spatial, mass and velocity distributions. The increased efficiency of dynamical friction and tidal disruption, in a triaxial galaxy, is so increased that their effects are not limited to the very inner regions of the parent galaxy. Moreover, there is no need to assume, for the GCS, a box-biased phase-space density such that the globular clusters are all on box orbits. In fact, Capuzzo Dolcetta (1993) showed that even in the case of an isotropic distribution function (hereafter DF), the evolution of the GCS is very similar to that of a box-biased DF. The reason for this is that the two DFs do not differ much in the region of the phase-space occupied by the majority of the clusters. In Section 2 we describe our model of a globular cluster system evolving due to dynamical friction and tidal disruption effects and we give a formula which permits to calculate the density profile of the GCS and its observable properties; in Section 3 we present and discuss the results. ", + "conclusions": "In this paper we have developed a model which allows to follow the evolution of the density distribution of a globular cluster system (GCS) in a triaxial galaxy under the influence of two effects: dynamical friction by field stars and tidal disruption by a massive central object. Both these effects are amplified by the triaxiality of the gravitational potential of the parent galaxy. The exact knowledge of the present-time density profiles permits to calculate the value of some of the relevant observables as a function of the assumptions made on the initial distribution and mass function of the GCSs. Actually, we find that the minor concentration of the GCSs, relatively to the distribution of halo stars, as observed in many galaxies, is an effect which depends strongly on the quality of the observations. In particular, a comparison of the core radius of the GCS with that ($r_{c0}$) of its parent galaxy star distribution might be misleading, since the estimate for the core radius of the GCS depends on the minimum radius at which clusters are sampled. For example, we have shown that, in the case of a GCS whose evolution is prevalently ruled by dynamical friction, as it is the case for a GCS made up mostly of heavy clusters, the final shape of the distribution displays a strong concentration of massive clusters in the very inner regions of the galaxy and a consequent lack of clusters in the outer regions. In this case, if the observations do not reach the inner regions (within $r_{c0}$), we effectively measure an erroneusly large core radius. Thus, an observed larger core radius is not a firm signature of the presence of a massive object at the centre of the galaxy. However, we found noticeable differences among the case of a steep IMF and that of a flat IMF. These differences, due to the influence of the half-mass density of the clusters as a function of the cluster mass, are such that: $i)$ for a steep IMF, made up mostly of light clusters, the prevailing effect is tidal disruption and there is a strong dependence of the evolved core radius on the mass of the central object (the value $r_c=10 r_{c0}$ is obtained with a nucleus of 2$\\cdot 10^8 M_\\odot$); $ii)$ for a flat IMF, made up mostly of heavy (and dense) clusters, the prevailing cause of depletion is dynamical friction and, consequently, the dependence on the mass of the central object is weak ($r_c= 10 r_{c0}$ for $2\\cdot 10^9 M_\\odot$). We have found that a parameter which is more reliable than the core radius to describe the GCS is the number of `lost' clusters outside some minimum radius (we choose $r_{min}=r_{c0}$). This is the difference between the initial and the observed number of clusters, under the assumption that the GCS was initially distributed as the parent halo light. Our calculations show that: for a flat IMF the percentage of `lost' clusters ranges from 40 per cent (with no massive central object) to 50 per cent (for a 10$^9 M_\\odot$ nucleus) of the initial total number; for a steep IMF it ranges from 3 per cent (no nucleus) to 80 per cent (10$^9 M_\\odot$ nucleus). We give two formulas which fit the fraction of `lost' clusters in the case of dynamical friction only and in the case of tidal disruption only, as a function of the cluster mass. When both the effects are working, the number of `lost' clusters may not be obtained by simply summing the fractions, confirming that an interaction among the two effects exists. To conclude, our calculations show that the present differences observed between the GCS and halo stars surface distributions {\\it can be explained by dynamical evolution of the GCS, under the influence of dynamical friction and tidal disruption}, even if the initial concentrations (core radii) were the same. Thus, to have a definite answer to the question \"are the observed differences between the star-bulge and the GCS density profiles just reflecting different initial conditions or are a consequence of evolution?\", it would be crucial to compare their kinematical properties. For example, the knowledge of the run with radius of the GCS (projected) velocity dispersion may help to understand if, like in our galaxy, there has been a selective depauperation of clusters on highly radial orbits." + }, + "9701/astro-ph9701060_arXiv.txt": { + "abstract": "{ Previous work reported a bar signature in color-selected IRAS variable stars. Here, we estimate the source density of these variables while consistently accounting for spatial incompleteness in data using a likelihood approach. The existence of the bar is confirmed with shoulder at $a\\approx4$ kpc, axis ratio $a:b=2.2$ -- $2.7$ and position angle of $19^{\\circ}\\pm1^{\\circ}$ degrees. The ratio of non-axisymmetric to axisymmetric components gives similar estimate for the bar size $a=3.3\\pm0.1$ kpc and position angle $\\phi_0= 24^{\\circ}\\pm2^{\\circ}$. We estimate a scale length $4.00\\pm0.55$ kpc for the IRAS variable population, suggesting that these stars represent the old disk population. We use this density reconstruction to estimate the optical depth to microlensing for the large-scale bar in the Galactic disk. We find an enhancement over an equivalent axisymmetric disk by 30\\% but still too small to account for the MACHO result. In addition, we predict a significant asymmetry at positive and negative longitudes along lines of sight through the end of the bar ($|l|\\approx30^\\circ$) with optical depths comparable to that in Baade's window. An infrared microlensing survey may be a sensitive tool for detecting or constraining structural asymmetries. More generally, this is a pilot study for Bayesian star count analyses. Bayesian approach allows the assessment of prior probabilities to the unknown parameters of the model; the resulting likelihood function is straightforwardly modified to incorporate all available data. However, this method requires the evaluation of multidimensional density functions over the data and optimization of the function over a parameter space. We address the resulting computational extremization problem with a hybrid use of a directed search algorithm which locates the global maximum and the conjugate gradient method which converges quickly near a likelihood maximum. Both methods are parallelizable and therefore of potential use with very large databases. } ", + "introduction": "\\label{sec:intro} Weinberg (1992, Paper I) identified color-selected variables in the IRAS Point Source Catalog (PSC) with AGB stars based on color consistency and the circumstantial sensitivity of the IRAS survey to long-period variables (cf. Harmon \\& Gilmore 1988). These were then used as rough standard candles to infer a large-scale asymmetry in the stellar distribution. The identification of IRAS variables with AGB stars was strengthened by an in-depth study of a bright subset (Allen, Kleinmann \\& Weinberg 1993). Carbon-selected AGB stars (carbon stars) have also proven to be effective tracers (see e.g. Metzger \\& Schechter 1994). Advantages of AGB tracers are reviewed in Weinberg (1994). In general, standard candle analyses have the advantage over flux or star count analyses in providing direct information about the three-dimensional structure of the Galaxy. However, uncertainties in their selection and intrinsic properties may bias any inference and, especially for the IRAS-selected sample, the census is incomplete. Paper I described an approach to large-scale Galactic structure using a star count analysis which allows the information to be reconstructed and possibly corrected in the observer's coordinate system before translating to a Galactocentric system. Unfortunately, this translation approach is only natural if the coverage is complete and suffered in application to the IRAS sample because of spatial gaps due to an incomplete second full-sky epoch. Here, we present the results of a different approach to the problem: the direct density estimation by maximum likelihood. A Bayesian density estimation has the advantage of directly incorporating selection effects and missing data. The number of ongoing surveys that bear on Galactic structure---SDSS, 2MASS, DENIS---which at various stages will have surveyed parts of the sky is a second motivation for this study; there is a need for a systematic method suited to inferential studies using possibly incomplete data from many wave bands. Recent analyses (e.g. Bahcall \\& Soneira 1980 in the optical; Wainscoat et al. 1992 in the infrared) have modeled the Galactic components with standard profiles and structural parameters chosen to provide a match to star count data. To explore the structural parameters themselves, we propose a Bayesian density estimation technique to treat data from scattered fields during the survey and to easily incorporate data from wave bands. Conceptually, this approach is midway between a classical inversion and modeling. The first part of the paper describes and characterizes the method. More specifically, \\S\\ref{sec:iras} reviews the IRAS selection procedure described in Paper I and motivates the approach. The new analysis based on statistical density estimation is presented in \\S\\ref{sec:bayes} and precisely defined in \\S\\ref{sec:likelihood}. The second part of the paper describes Monte-Carlo tests and the results of applying the method to the IRAS data (\\S\\ref{sec:results}). We conclude in \\S\\ref{sec:summary} with a summary and discussion. ", + "conclusions": "\\label{sec:summary} This paper explores a model-independent Bayesian estimation of the stellar density from star counts, rigorously accounting for incomplete data. The general approach can incorporate multiple colors and even different databases. The usual high dimensionality and topological complexity of the posterior distribution, however, complicates both optimization algorithms and subsequent moment analyses. We propose here a hybrid downhill plus directed-search Monte Carlo algorithm; the former speeds convergence and the latter facilitates the location of the global extremum. Other similar and potentially more efficient techniques which can bypass the extremization step altogether (such as general Markov Chain Monte Carlo) are worth careful consideration. Application of the technique to the variability-selected sample described in Weinberg (1992), assumed to be AGB stars, confirms the presence of a strong non-axisymmetric feature in the first Galactic quadrant. By imposing bisymmetry on the source density, clear signature of a bar is obtained. The size and shape of density isophotes suggests a bar semi-major axis of approximately 4 kpc and position angle of $\\phi_0 = 18^\\circ \\pm 2^\\circ$ at the outer edge of the bar. The analysis of the scale length for the AGB candidate distribution gives $r_0=4.00\\pm0.55$ kpc, indicating that these objects are part of the old disk population. Finally, we use our estimate for non-axisymmetric Galactic disk to explore the dependence of optical depth to gravitational microlensing by bulge and disk stars. The disk bar does enhance the optical depth $\\tau$ towards Baade's window by roughly 30\\% but the overall value is still roughly a factor of two below the MACHO result $\\tau = 3.9^{+1.8}_{-1.2} \\times 10^{-6}$. Of interest for future microlensing surveys is the finding that our inferred large-scale bar will produce a significant asymmetry in $\\tau$ at positive and negative longitudes beyond the bulge. The peak asymmetry for our model occurs at $|l|=30^\\circ$ and at $b=0$ we predict similar values of $\\tau$ to the Baade's window field. Such a survey might best be carried out in the infrared to take advantage of the low interstellar extinction and colors of the late-type giants. At $|l|\\gtrsim30^\\circ$, confusion should not be a limitation at $b=0^\\circ$." + }, + "9701/astro-ph9701195_arXiv.txt": { + "abstract": "We present a new high-resolution $N$-body algorithm for cosmological simulations. The algorithm employs a traditional particle-mesh technique on a cubic grid and successive multilevel relaxations on the finer meshes, introduced recursively in a {\\em fully adaptive} manner in the regions where the density exceeds a predefined threshold. The mesh is generated to effectively match an {\\em arbitrary} geometry of the underlying density field -- a property particularly important for cosmological simulations. In a simulation the mesh structure is not created at every time step but is properly adjusted to the evolving particle distribution. The algorithm is fast and effectively parallel: the gravitational relaxation solver is approximately half as fast as the fast Fourier transform solver on the same number of mesh cells. The required CPU time scales with the number of cells, $N_c$, as $\\sim O(N_c)$. The code allows us to improve considerably the spatial resolution of the particle-mesh code without loss in mass resolution. We present a detailed description of the methodology, implementation, and tests of the code. We further use the code to study the structure of dark matter halos in high-resolution ($\\sim 2h^{-1}$ kpc) simulations of standard CDM ($\\Omega=1$, $h=0.5$, $\\sigma_8=0.63$) and $\\Lambda$CDM ($\\Omega_{\\Lambda}=1-\\Omega_0=0.7$, $h=0.7$, $\\sigma_8=1.0$) models. We find that halo density profiles in both CDM and $\\Lambda$CDM models are well fitted by the analytical model presented recently by Navarro et al., which predicts a singular [$\\rho(r)\\propto r^{-1}$] behavior of the halo density profiles at small radii. We therefore conclude that halos formed in the $\\Lambda$CDM model have structure similar to CDM halos and thus cannot explain the dynamics of the central parts of dwarf spiral galaxies, as inferred from the galaxies' rotation curves. ", + "introduction": " ", + "conclusions": "" + }, + "9701/astro-ph9701148_arXiv.txt": { + "abstract": "The mean intracluster gas fraction of \\xray clusters within their hydrostatic regions is derived from recent observational compilations of David, Jones \\& Forman and White \\& Fabian. At radii encompassing a mean density 500 times the critical value, the individual sample bi-weight means are moderately ($2.4 \\sigma$) discrepant; revising binding masses with a virial relation calibrated by numerical simulations removes the discrepancy and results in a combined sample mean and standard error $\\fbargas(\\rfiveh) = (0.060 \\pm 0.003) \\h^{-3/2}$. For hierarchical clustering models with an extreme physical assumption to maximize cluster gas content, this value constrains the universal ratio of total, clustered to baryonic mass $\\Omega_m/\\Omega_b \\le 23.1 \\h^{3/2}$; combining with the primordial nucleosynthesis upper limit on $\\Omega_b$ results in $\\Omega_m \\h^{1/2} < 0.60$. A less conservative, physically plausible approach based on low D/H inferences from quasar absorption spectra and accounting for baryons within cluster galaxies yields an estimate $$ \\Omega_m \\h^{2/3} = 0.28 \\pm 0.07 $$ with sources of systematic error involved in the derivation providing approximately $35\\%$ uncertainty. Additional effects which could provide consistency with the Einstein--deSitter case $\\Omega_m \\se 1$ are presented, and their observable signatures discussed. ", + "introduction": "Clusters of galaxies provide a number of interesting cosmological diagnostics. In particular, the relative amount of baryons and dark matter within their hydrostatic regions provides a measure of the cosmic mix of these components, \\ie, a measure of the ratio of density parameters $\\Omega_m/\\Omega_b$ in the Friedmann--Lema\\^{\\i}tre world model\\footnote{In this paper, $\\Omega_m$ refers to the contribution of all clustered matter (including baryons) to the stress--energy density.} Employing this measure in practice requires accurate observational data along with estimates of possible systematic biases. Likely sources of bias include systematic errors in component mass estimates and deviation of the local cluster ratio of baryonic--to--total mass arising, for example, from the different dynamical histories of the two components. Within the context of hierarchical clustering scenarios, these effects have been calibrated by numerical simulations of cluster formation. The results, discussed in detail below, indicate that the magnitude of these effects are small --- a few tens of percent or less --- in observationally accessible regions of clusters. If our current description of cluster formation dynamics is physically accurate, then determination of the cosmic baryon fraction from \\xray cluster data is straightforward. The baryonic component of the richest clusters is dominated by the \\xray emitting intracluster gas rather than the mass associated with the optical light of the cluster galaxies (Forman \\& Jones 1982; Sarazin 1986; Mushotzky 1994). Taking Coma as an example, White \\etal (1993) estimate the ratio of gas to galaxy mass to be $M_{gas}/M_{gal} = (5.5 \\pm 1.5) \\h^{-3/2}$ (with $h \\! = {\\rm H}_0/100 \\kmsmpc$). This number is typical of the values seen in larger samples for clusters comparable in temperature to Coma (David \\etal 1990; Arnaud \\etal 1992). The data also indicate a dependence of the ratio $M_{gas}/M_{gal}$ on cluster temperature, with poorer clusters and groups possessing less gas per massive galaxy than their larger counterparts. The gas mass thus provides not only a formal lower limit to the total baryon cluster content, but a fairly accurate estimate of the baryon mass in rich clusters if the dark matter is assumed non--baryonic. Recent compilations of \\xray cluster data by White \\& Fabian (1995) and David, Jones \\& Forman (1995) provide gas and total mass estimates for samples of moderate size. In this paper, I use these data, along with guidance from numerical simulations, to estimate the sample mean gas mass fraction $\\fbargas$ within a characteristic radius defining the hydrostatic boundary of clusters. Motivations for the specific approach adopted here are provided in \\S2 and the observational data are analysed in \\S3. Implied constraints on the universal baryon fraction and the value of $\\Omega_m$ are provided in \\S4, including a thorough discussion of systematic errors. ", + "conclusions": "The gas fraction in clusters of galaxies provides information on the cosmic baryon mass fraction. Dynamical biases of clusters' baryon content can be minimized by measuring masses near the virial radius, where gas dynamic experiments show equilibrium is valid and segregation processes inefficient. At $\\rfiveh$, the radius where the mean interior density is a factor 500 times the critical value, the bi--weight mean gas fraction of the combined, revised WF and DJF \\xray cluster samples is $0.060 \\pm 0.003 h^{-3/2}$. This value, when combined with our current understanding of cluster formation history and limits on $\\Omega_b$ from primordial nucleosynthesis, strongly favors $\\Omega_m \\h^{2/3} \\sim 0.3$ and rules out the possibility of $\\Omega_m \\se 1$ with high statistical significance, unless the Hubble constant is very low $h \\! < \\! 0.4$. These conclusions reinforce, at greater statistical significance, earlier work based on different methods and data than those used here (\\eg, Henriksen \\& Mamon 1994; Steigman \\& Felten 1995; Lubin \\etal 1996). How serious is the case against $\\Omega_m \\se 1$? A die--hard Einstein--deSitter advocate could espouse all the elements of the analysis leading to the upper limit in eq'n~(\\ref{omega_upper}) and claim only a small ($30\\%$ for $h \\se 0.7$) systematic effect is missing. However, this approach accounts for neither the hot, \\xray emitting gas nor the galaxies in clusters, and one might well be suspicious of an argument which ignores these two principal, observable components. The best estimate approach leaves one short by a factor $\\gta 3$. It is possible that this gap is plugged not by a single large effect, but by several effects acting in concert, a situation reminiscent of the so--called ``$\\beta$--discrepancy'' in clusters (Evrard 1990b; Lubin \\& Bahcall 1995). There remains the possibility that a new element of \\xray cluster physics is simply missing from the present picture. Observational constraints on known sources of systematic error should be vigorously pursued, along with more sophisticated theoretical modeling of intracluster plasma dynamics and thermodynamics. The arguments presented here are independent of the value of the cosmological constant. If one favors a spatially flat universe, as motivated by simple models of inflation, obtained through a non-zero cosmological constant $\\Lambda$, then the limits on the clustered matter component predict values of the required vacuum energy density $\\Omega_\\Lambda \\equiv \\Lambda/3\\Ho^2 = 1 - \\Omega_m$. For example, for $h=0.7$, the upper limit on $\\Omega_m$ requires $\\Omega_\\Lambda \\ge 0.28$ while the best estimate approach gives $\\Omega_\\Lambda = 0.64 \\pm 0.09$. The latter is consistent with the limit $\\Omega_\\Lambda < 0.66$ derived from gravitational lensing arguments (Kochanek 1996) while marginally inconsistent with the $\\Omega_\\Lambda < 0.51$ inferred recently by Perlemutter \\etal (1996) from the magnitude--redshiftt relation for Type-Ia supernovae." + }, + "9701/astro-ph9701181_arXiv.txt": { + "abstract": "The production of \\al by explosions of classical novae has been computed by means of a hydrodynamic code that follows both the accretion and the explosion stages. A special emphasis has been put on the analysis of the influence of the initial abundances of the accreted envelope, as well as on the nuclear reaction rates involved. With the most recent values of chemical composition and reaction rates available, \\al production is lowered with respect to previous computations. According to our results, the final contribution of novae to the galactic \\al is at most 0.4~\\msun, which is a small part of the estimated \\al in the Galaxy derived from COMPTEL observations of the 1809 keV emission. ", + "introduction": "The discovery of \\al in the interstellar medium by the HEAO-3 satellite, through the detection of the 1809 keV $\\gamma$-ray line (\\cite{Mah82}, \\cite{Mah84}), has raised the interest on potential $\\gamma$-ray diagnostics of nucleosynthesis in several scenarios. This $\\gamma$-ray line is produced by the decay ($\\tau\\sim 10^6$ years) from the ground state of \\al to the first excited state of \\mggg, which de-excites to its ground state level by emitting a 1809 keV photon. This detection has been confirmed by other space missions like the SMM $\\gamma$-ray spectrometer (\\cite{Sha85}) or several balloon-borne experiments. Recent measurements made with the COMPTEL instrument on-board the Compton Gamma-Ray Observatory (CGRO) have provided a map of the 1809 keV emission in the Galaxy. This map shows an extended diffuse emission along the galactic plane, with a peculiar large scale asymmetry about the galactic center and a clumpy structure with several noticeable hot spots (\\cite{Die95}). The first interpretations of this map have suggested that novae and low--mass stars cannot be the major contributors to the 1809 keV emission, because their low individual yields and their high frequency should provide a smooth distribution, in contradiction with the irregular appearance of the emission (see \\cite{PD96} for a review). The general model suggested considers a two-component origin of the emission: a global component following the spiral pattern of the Galaxy (a presumed site of massive star formation) on which several localized regions of intense activity, such as the Vela region, are superimposed (\\cite{CGD95}). However, observations of four nearby supernova remnants with COMPTEL do not provide evidence for the 1809 keV emission, although the uncertainties in the distances do not allow to put severe constraints on \\al production by supernovae (\\cite{Kno96a}). Finally, recent attempts to understand the 1809 keV map of the sky have analyzed in more detail its correlation with the spiral structure (\\cite{Kno96b}). They derive a total \\al mass of 2.5 \\msun, from which at least 0.7 \\msun can be attributed to massive stars, stressing that they cannot exclude that a large fraction of \\al is produced by novae or low--mass AGB stars. Therefore, with this still unclear panorama, it is worth studying the role played by classical novae in the synthesis of galactic \\al. As pointed out by \\cite{WF80}, \\al production requires moderate peak temperatures, of the order of $\\rm{T_{peak}} \\leq ~2 \\power{8}$ K, and a fast decline from maximum temperature, conditions that are commonly achieved in nova outbursts. In the early 80s, one-zone model calculations of explosive hydrogen burning nucleosynthesis by \\cite{HT82} and \\cite{Wie86} suggested that classical novae might produce sufficient amounts of \\al to account for some of the observed meteoritic anomalies but would not represent major galactic sources. But these calculations used only solar or CNO-enhanced envelopes. New one-zone model calculations by \\cite{WT90} and \\cite{Nof91}, on the basis of ONeMg white dwarf stars, produced large amounts of long-lived radioactive nuclei, such as \\sod and \\al, concluding that these novae might represent important (though not dominant) sources of the galactic \\al. However, according to \\cite{Nof91}, no \\al production results for models with T$_{\\rm {peak}} \\geq~2.7 \\power{8}$ K. Recent calculations have refuted this result (\\cite{Pol95}, \\cite{Sta93}), demonstrating the crucial role played by convection in order to carry some \\al to the outer, cooler layers of the envelope, where its destruction through proton captures is prevented. But these fully hydrodynamical calculations involving ONeMg white dwarfs assume a composition for the white dwarf core (which is partially mixed with the envelope) based on old calculations of C-burning nucleosynthesis by \\cite{AT69} that need to be updated. Since the resulting nucleosynthesis is very sensitive to the envelope's composition, it is important to adopt a more realistic one. On the other hand, there are large uncertainties on some crucial reactions involving \\al production. Thus, it is important to analyse the influence of the different prescriptions available on the final \\al yields in the ejecta of classical novae. With this aim we have computed a series of hydrodynamic models of nova outbursts for several white dwarf masses and accretion rates, adopting updated initial chemical compositions and recent prescriptions for the crucial reaction rates. ", + "conclusions": "The production of \\al by classical novae is very sensitive to the initial composition of the envelope and to the nuclear reaction rates adopted. ONe novae are more important \\al producers than CO ones, because seed nuclei for the NeNa and MgAl cycles are almost absent in the latter ones. For the same reason, the amount of \\al synthesized in ONe novae depends on the initial composition of the white dwarf core. When using the recently available chemical abundance profiles by \\cite{Rit96}, a lower \\al production is obtained. Furthermore, some improvements in the nuclear reaction rates since \\cite{CF88} also lead to a lower \\al production. The amount of \\al injected into the interstellar medium by ONe nova events decreases as the mass of the underlying white dwarf increases. Hence, low-mass white dwarfs are most likely candidates for \\al production. Contribution from low-mass white dwarfs is favored both by the higher \\al production and the higher ejected mass. But white dwarfs of masses lower than $\\sim$1.1 \\msun are expected to be CO white dwarfs, which are unable to produce important quantities of \\al (see table 1). Also, the observations of some Ne-novae, such as QU Vulpecula 1984, indicate high ejected masses ($\\sim 10^{-3}$ \\msun), unobtainable if a massive ONe nova is the responsible of the explosion. This has led some authors to propose a new scenario (\\cite{Sha94}, \\cite{SP94}): low-mass CO white dwarfs, undergoing episodic accretion phases at high rates, experience large metal enrichments from the ashes of He-burning. The subsequent nova explosion, correlated with a phase of a lower accretion rate, produces at the same time high \\al abundances and high ejected masses. But these models require a fine tuning of some parameters relative to the cataclysmic variable, such as the mass accretion rate. A crude estimate of the total production of \\al by novae can be computed by means of the expression from \\cite{WT90}: $$\\rm M({}^{26}Al) = 0.4 { X({}^{26}Al) \\over 2 \\power{-3} } {frac(Ne) \\over 0.25} {M_{ej} (M_\\odot) \\over 2 \\power{-5} } {R_{nova} (events~yr^{-1}) \\over 40}$$ \\noindent where X(\\al) represents the mean mass fraction of \\al in the ejecta, frac(Ne) the fraction of Ne-novae over the total number of classical nova outbursts (between 0.25 and 0.53, but typically $\\sim 1/3$, see \\cite{TL86}, \\cite{LT94}), $\\rm M_{ej}$ the amount of mass ejected in an outburst, and $\\rm R_{nova}$ the nova rate ($\\sim 46 \\rm yr^{-1}$, \\cite{HF87} or $\\sim 20 \\rm yr^{-1}$, \\cite{DL94}). Adopting these estimates of R$_{\\rm{nova}}$ and frac(Ne) and our most favorable ONe nova case, for which $\\rm M_{ej}$~(\\al) = 1.7 \\power{-8} \\msun, we derive a maximum contribution of classical nova outbursts to the galactic \\al in the range 0.1 to 0.4 \\msun. Although higher abundances of \\al can be obtained for models with higher enrichment from core material (i.e., 75\\%), these may be considered as very extreme cases, in view of the typical metallicities observed in the ejecta of classical nova outbursts. In summary, the contribution of novae to galactic \\al is small as compared with the one required to explain COMPTEL measurements (between 1 and 3 \\msun). This is in agreement with the most accepted hypothesis of young progenitors as sources of galactic \\al. We want to stress that the two aspects that are crucial for the final \\al yields are far from being understood; these are the process of mixing between the core and the envelope, which determines the initial chemical profile, and the exact rates of the nuclear reactions involved in \\al synthesis. This latter topic will surely be improved with the new compilation of nuclear reaction rates (\\cite{NACRE})." + }, + "9701/astro-ph9701238_arXiv.txt": { + "abstract": "The ohmic decay of magnetic fields confined within the crust of neutron stars is considered by incorporating the effect of space-time curvature produced by the intense gravitational field of the star. It is shown that general relativistic effect reduces the magnetic field decay rate substaintially and specially at the late time of the evolution the decay rate decreases by several orders of magnitude when compared with the case without the relativistic effect. ", + "introduction": "The magnetic field evolution in neutron stars has been a subject of much discussion over the years both in the observational and in the theoretical context. Calculations of ohmic decay of dipolar magnetic fields were performed by Sang \\& Chanmugam (1987) who demonstrated that the field does not decay exponentially. The reviews by Lamb (1991), Chanmugam (1992) and Phinney \\& Kulkarni (1994) provide the present understanding on the decay of magnetic fields in isolated neutron stars. Haensel, Urpin \\& Yakovlev (1990) pointed out to the possibility that magnetic fields in the core could decay rapidly by ambipolar diffusion. The studies of magnetic field configurations in which the field vanishes in the stellar core (\\cite{SC87}; \\cite{CS89}; \\cite{UM92}; \\cite{UV93}) show that because of the relatively low electrical conductivity of crustal matter, the decay times may be short enough to be of observational interest if the impurity concentration is high and the field is initially confined to a small part of the crust. On the other hand, considering field configurations that do not vanish in the core, Pethick \\& Sahrling (1995) showed that the shortest possible decay time is about two orders of magnitude longer than the characteristic time-scale for decay of configurations in which the magnetic field vanishes in the core. In all the investigations, an important feature of neutron stars, the space-time curvature produced by the intense gravitational field has not been taken into account although it is well-known that the space-time curvature exterior and interior to neutron stars can significantly alter the electromagnetic field (\\cite{WS83}; \\cite{S95}; \\cite{BO95}). In this letter it is demonstrated, by adopting a simplified model, that the decay rate decreases significantly when the general relativistic effect is taken into consideration. It is, however, worth mentioning that due to lack of proper understanding on the initial configuration of the magnetic field and of the impurity parameter which plays crucial role in determining the electrical conductivity, no model presented so far can be attributed to represent the quantitative feature of the actual situation. Nevertheless, all these investigations are important from the qualitative point of view and provide significant insight on the decay of the magnetic field. The scope of the present work although is also limited to an idealistic situation but it is sufficient to demonstrate the important role played by curved space-time at the crust of neutron stars in governing the decay of magnetic fields. The present result indicates that the space-time curvature produced by the intense gravitational field at the crust can lead to characteristic decay-times much longer than the existing estimations. ", + "conclusions": "The important message which is conveyed by the present calculations is that whatever be the electrical conductivity of the crustal material, high or low, irrespective of the question whether the magnetic field vanishes at the core or not and whatever be the impurity content of the neutron star crust, the decay time of the magnetic field is lengthened by the intense gravitational field that certainly exists inside the star. Sang \\& Chanmugam (1987) showed that the decay is not exponential while Urpin \\& Muslimov (1992) pointed out that even if the magnetic field is initially absent in the core, diffusion of the field into a highly conducting core would retard the surface field decay. Pethic and Sahrling (1995) suggested that if long decay times were established observationally, these could not necessarily imply as the evidence for matter in the stellar core having a high conductivity. Irrespective of all the uncertainties that still exist and require further theoretical investigations as well as observational evidences the present result can atleast provide a firm understanding that general relativity is certainly responsible if the decay time is indeed very long. Therefore, the present demonstration is important in the sense that it establishes a concrete restrictions on the theoretically possible ways of obtaining short ohmic decay times for magnetic fields in neutron stars. Further, the present results provide an interesting feature that the more is the compactness of the neutron star the longer is the decay times and hence the general relativistic effects on the decay of magnetic field could be a possible tool for restricting the equation of state of matter inside the star which determines the compactness of the neutron star. The crucial lesson that the present results provide is that, in addition to the calculations of the conductivity with better estimation of the impurity content, the incorporation of the effect of superfluidity and superconductivity in the core and other physical effects such as Hall drift, full consideration of the general relativistic effects must be given in order to make more realistic estimates of decay times of the magnetic field." + }, + "9701/astro-ph9701074_arXiv.txt": { + "abstract": "After briefly summarizing the main tenets of unified schemes of Active Galactic Nuclei, I review some recent results in the field of unification of radio-loud sources, both for the low-luminosity (BL Lacs and Fanaroff-Riley type I radio galaxies) and high-luminosity (radio quasars and Fanaroff-Riley type II radio galaxies) populations. ", + "introduction": "It now seems well established that the appearance of Active Galactic Nuclei (AGN) depends strongly on orientation. Classes of apparently different AGN might actually be intrinsically similar, only seen at different angles with respect to the line of sight. The basic idea, based on a variety of observations and summarized in Figure 1 of Urry \\& Padovani (1995), is that emission in the inner parts of AGN is highly anisotropic. The current paradigm for AGN includes a central engine, surrounded by an accretion disk and by fast-moving clouds, probably under the influence of the strong gravitational field, emitting Doppler-broadened lines. More distant clouds emit narrower lines. Absorbing material in some flattened configuration (usually idealized as a torus) obscures the central parts, so that for transverse lines of sight only the narrow-line emitting clouds are seen (narrow-lined or Type 2 AGN), whereas the near-IR to soft-X-ray nuclear continuum and broad-lines are visible only when viewed face-on (broad-lined or Type 1 AGN). In radio-loud objects we have the additional presence of a relativistic jet, roughly perpendicular to the disk, which produces strong anisotropy and amplification of the continuum emission (``relativistic beaming''). In general, different components are dominant at different wavelengths. Namely, the jet dominates at radio and $\\gamma$-ray frequencies (although it does contribute to the emission in other bands as well), the accretion disk is thought to be a strong optical/UV/soft X-ray emitter, while the absorbing material will emit predominantly in the IR. It then follows that a proper understanding of AGN will only come through multifrequency studies (e.g., Padovani 1997). This axisymmetric model of AGN implies widely different observational properties (and therefore classifications) at different aspect angles. Hence the need for ``Unified Schemes'' which look at intrinsic, isotropic properties, to unify fundamentally identical (but apparently different) classes of AGN. Seyfert 2 galaxies have been ``unified'' with Seyfert 1 galaxies (see Antonucci 1993, and references therein, and Granato, these Proceedings, for more recent results). As regards the radio-loud population, Fanaroff-Riley type I (FR I: i.e., low-power) radio galaxies have been unified with BL Lacertae objects, a class of AGN characterized by very weak emission lines, while Fanaroff-Riley type II (FR II: i.e., high-power) radio galaxies have been unified with radio quasars. In the latter case, flat-spectrum radio quasars (FSRQ) are thought to be oriented at relatively small angles w.r.t. to the line of sight ($\\theta \\simlt 15^{\\circ}$), while steep-spectrum radio quasars (SSRQ) should be at angles intermediate between those of FSRQ and FR II radio galaxies. This paper is {\\it not} meant to be a review of unified schemes: the interested reader might consult, for example, the reviews by Antonucci (1993) and Urry \\& Padovani (1995). My aim is to discuss some of the very recent results in the field of unified schemes for radio-loud AGN which have appeared in the literature in the last year or so (and therefore were not included in Urry \\& Padovani 1995). ", + "conclusions": "I have listed here some of the open questions/hot topics on which work is in progress. These will be hopefully discussed at the next Italian meeting on AGN, to be held in 1998! \\begin{itemize} \\item Obscuration in FR Is. The presence of broad lines in at least some BL Lacs (Sect. 2.1) indicates that some obscuring material must be present in FR Is (as well as in FR IIs). A SAX core program will address this issue. Also, ISO observations of FR Is should provide further constraints. \\item BL Lac (and FR I) environment. As stressed in Section 2.2, it might well be that BL Lacs avoid the richest clusters. More work in this field is needed, as the samples used to study the environment of BL Lacs are still relatively small. Various ongoing projects will address this point in more detail. Comparable studies on the environment of FR I radio galaxies using sizeable and well defined samples should also be performed. \\item HBL/LBL dichotomy. As mentioned in Section 2.3, we need deeper and larger BL Lac samples to address this point. Amongst the various groups working on this, both Perlman, Padovani, Giommi, \\ea (1997) and Caccianiga, Maccacaro, Wolter, \\ea (these Proceedings) are cross-correlating X-ray and radio catalogues to select new BL Lac (and FSRQ) samples. \\item (Spectro)Polarimetry of complete samples of radio-galaxies. As mentioned in Section 3.2, spectropolarimetry allows one to see emission from the nuclear region, particularly broad lines, even in Type 2 sources. So far, however, spectropolarimetric studies have concentrated on individual, sometimes peculiar, sources. We would need a spectropolarimetric study of a complete, well-defined sample, to quantify how common reflected broad lines are, for example. As far as I know, nobody is doing this, but I hope that I am misinformed on this! Tadhunter \\ea and di Serego Alighieri et al., however, are carrying out a polarimetric study of a complete sample of radio sources, which should still give interesting, quantitative information. \\item VLBI studies of complete sample of radio sources. These are important to study, among other properties, the distribution of superluminal speeds to estimate beaming parameters and constrain unification. Various surveys are in progress, for example the Caltech-Jodrell Bank survey of strong ($0.7 < f_{\\rm 6cm} < 1.3$ Jy) radio sources (e.g., Polatidis \\ea 1995) or the survey of low-power radio galaxies by the Bologna group (e.g., Venturi \\ea 1995). \\end{itemize} In conclusion, the unification of radio-loud sources is a very active field, with interesting results appearing even on a relatively short time scale. In 1995 alone, about 65 papers on beaming and unification of radio sources have been published (51 in refereed journals), an all-time record judging from the data relative to the last ten years (source: Astrophysics Data System). Various modifications to our relatively simple and idealized unification picture seem to be in order: not all FR I radio galaxies might host BL Lacs, after all, but probably only those in relatively poor clusters; obscuring material might be present in FR I radio galaxies as well, and not only in their higher-luminosity counterparts; not unexpectedly, radio sources do evolve in time, growing in size and dimming in power, which might help solving some inconsistencies in their linear sizes relative to those of radio galaxies; infrared broad lines are indeed detected in about half of the low-redshift narrow line radio galaxies; but we still do not know if BL Lacs of the LBL type are intrinsically more numerous than BL Lacs of the HBL type! Finally, although so far unified schemes for radio-loud and radio-quiet AGN have been considered separately, there might be some connection between the two. Falcke, Sherwood \\& Patnaik (1996) have recently suggested that relativistic boosting, normally associated with radio-loud AGN, might be present in the radio cores of some radio-quiet sources as well. This idea, with its far-reaching implications, might open up new research paths which could help us to solve the long-standing problem of the radio-loud/radio-quiet dichotomy of the AGN population." + }, + "9701/astro-ph9701132_arXiv.txt": { + "abstract": "We report the detection of \\ion{N}{5} 1239\\AA\\ transition region emission in {\\it HST}/GHRS spectra of the A7 V stars, $\\alpha$ Aql and $\\alpha$ Cep. Our observations provide the first direct evidence of $1-3 \\times 10^5$ K material in the atmospheres of normal A-type stars. For both stars, and for the mid-A--type star $\\tau^3$~Eri, we also report the detection of chromospheric emission in the \\ion{Si}{3} 1206\\AA\\ line. At a \\bv color of 0.16 and an effective temperature of $\\sim8200$ K, $\\tau^3$~Eri becomes the hottest main sequence star known to have a chromosphere and thus an outer convection zone. We see no firm evidence that the \\ion{Si}{3} line surface fluxes of the A stars are any lower than those of moderately active, solar-type, G and K stars. This contrasts sharply with their coronal X-ray emission, which is $>100$ times weaker than that of the later-type stars. Given the strength of the \\ion{N}{5} emission observed here, it now appears unlikely that the X-ray faintness of the A stars is due to their forming very cool, $\\leq$1 MK coronae. An alternative explanation in terms of mass loss in coronal winds remains a possibility, though we conclude from moderate resolution spectra of the \\ion{Si}{3} lines that such winds, if they exist, do not penetrate into the chromospheric \\ion{Si}{3}--forming layers of the star, since the profiles of these lines are {\\it not} blueshifted, and may well be redshifted with respect to the star. ", + "introduction": "In this paper we address a long-standing question in stellar structure as to the locus for the onset of subsurface convection zones in main sequence stars. Thirty years ago, the changeover from radiative to convective envelopes was presumed to occur in the close vicinity of spectral type F5 (e.g., Wilson 1966, Kraft 1967), an idea that was supported by optical observations of the time, which showed an absence of chromospheric \\ion{Ca}{2} emission reversals in the spectra of earlier stars and a steep falloff in the axial rotation rates of later stars (the direct result, it was reasoned, of magnetic braking and angular momentum loss due to the onset of coronal winds). In later years, however, improved stellar structure models as well as X-ray and ultraviolet (UV) observations from space have shown that coronae, chromospheres, and thus convection zones, must be present in stars that are located well above the F5 boundary line. UV spectra from the {\\it International Ultraviolet Explorer (IUE)} telescope, for example, trace chromospheric emission in \\lya\\ and in the 1335 \\AA\\ lines of \\ion{C}{2} along the main sequence into the late-A stars, near \\bv = 0.20 (Simon \\& Landsman 1991; Marilli et al. 1992; Landsman \\& Simon 1993), while a number of A stars have been detected in X rays by the {\\it Einstein Observatory} or by {\\it ROSAT} (Schmitt et al. 1985; Simon, Drake, \\& Kim 1995), including a handful of A0 stars with $\\bv \\approx 0.0$. The X-rays of the cooler A-type stars are believed to be coronal in origin. Those of the hotter stars are thought to come not from the A stars themselves, but from hidden late-type binary companions or neighboring stars that lie within the X-ray beam. Only rarely can this ambiguity be resolved (e.g, Schmitt \\& K\\\"urster 1993). UV spectroscopy largely avoids this difficulty with the imagery, since the high temperature UV lines of a bright A star are expected to outshine those of virtually any low mass companion. On the other hand, the detection of chromospheric emission against the bright photosphere of an A star poses challenges of its own, demanding moderate-to-high spectral resolution, minimal scattered light, and exceptional signal-to-noise ratio. Such requirements were often beyond the capabilities of {\\it IUE}, but can now be met routinely by instruments aboard {\\it HST}, such as the Goddard High Resolution Spectrograph (GHRS). In previous work with the GHRS (Simon, Landsman, \\& Gilliland 1994, hereafter SLG), we obtained high quality moderate resolution spectra of the \\ion{C}{2} 1335 \\AA\\ lines for 8 mid- to late-A stars. Only in the case of the A7 V star, $\\alpha$~Aql (Altair), did we find evidence for chromospheric emission. We chose the \\ion{C}{2} lines for this earlier study because they are among the brightest features in the UV spectra of G--M stars; moreover, the photosphere of an A star is much fainter at 1335 \\AA\\ than it is at longer wavelengths, where the increasing brightness of the background can render even strong lines, such as \\ion{C}{4} 1550 \\AA, completely invisible (Simon \\& Landsman 1991). The strongest A star chromospheric line accessible to the GHRS is \\lya. This line suffers the dual disadvantage of contamination by scattered light in the Earth's atmosphere and attenuation by interstellar \\ion{H}{1} along the line of sight to the star. Nonetheless, even at low resolution and low S/N, {\\it IUE} spectra have yielded some detections among the late-A and early-F stars (Landsman \\& Simon 1993). Accordingly, with this {\\it IUE} work in mind, we have initiated a program of GHRS \\lya\\ spectroscopy for a modest sample of bright A stars. Our goal was to establish more stringent constraints on the locus for the onset of main sequence convection zones. In this article, we report the initial results of our observations, which include the first detection of chromospheric emission in a mid-A type star, as well as the first detections of $\\sim10^5$ K \\ion{N}{5} 1239\\AA\\ transition region emission in the spectra of two late A-type stars. ", + "conclusions": "To gauge the intrinsic activity of the A stars, we use the Barnes-Evans (1976) relation to convert the apparent fluxes in Table 1, $f_{\\lambda}$, to surface fluxes at the star, $F_{\\lambda}$. The results are in Table 2, along with comparable numbers for X-ray emission. For comparison, we also tabulate surface fluxes for representative G and K stars. These values derive mainly from GHRS spectra (e.g., Ayres et al. 1997) or, in some cases, from compilations of {\\it IUE} and {\\it ROSAT} data that have appeared in the literature (e.g., Ayres et al. 1995). In terms of their chromospheric or transition region fluxes, the A stars resemble moderately active late-type dwarfs or giants (e.g., $\\epsilon$ Eri or $\\beta$ Cet), but they are (to no surprise) much less active than the rapidly-rotating Hertzsprung Gap giants (e.g., 31 Com) and the tidally locked RS CVn binaries (e.g., V711 Tau). In coronal X rays, the A stars are comparable to middle aged stars like the Sun and $\\alpha^1$ Cen. They are substantially less active than prototypical, rapidly rotating, young dwarf stars like $\\chi^1$ Ori. Nor are they as active as any of the late-type giants. These trends are more clearly illustrated in the H--R bubble diagram shown in Figure 3. In this figure, we scale the diameter of the circle for each star to the common logarithm of the normalized flux value, $F_{\\lambda}/\\sigma T_{\\rm eff}^4$. Turning to the top panel, we find no evidence in the \\ion{Si}{3} line for a decline in chromospheric activity along the main sequence, even among the middle A stars (that is, for effective temperatures as high as 8200 K) if one includes the uncertain measurement of $\\tau^3$~Eri. By contrast, the coronal X-ray fluxes in the lower panel of the figure show an enormous decline from one side of the diagram to the other, with the main sequence and more evolved early type stars to the left being considerably fainter than the late-type stars on the right. Attention was drawn to the lower coronal/chromospheric flux ratios among the early type stars by Simon \\& Drake (1989), who ascribed the effect to mass loss via coronal winds; a more recent discussion by Ayres et al. (1997) concludes that the origin of these ``X-ray deficits''\\/ may be related to the possible existence of very long coronal loops, but otherwise the trend remains unexplained. At intermediate, transition region temperatures, a similar trend may be appearing in the \\ion{N}{5}/\\ion{Si}{3} and \\ion{N}{5}/\\ion{C}{2} ratios, each being a factor of $2-3\\times$ smaller among the A stars than among the G--K stars. However, given the lack of truly adequate numbers of high quality spectra of \\ion{N}{5} for stars of all MK classes, this possibility requires further study with the {\\it HST}. Continuing the tendency toward a broad dispersion in line strengths among the early F stars that was noted in our previous work (Simon \\& Landsman 1991; Landsman \\& Simon 1993), the surface fluxes of the A stars in Table 2 indicate that Altair is at least a factor of 2 more active than $\\alpha$ Cep over the entire range of temperatures from chromospheric to coronal, despite the otherwise close similarities of these two stars ($\\alpha$ Cep has the slightly lower surface gravity). Our spectra also largely put to rest the notion that the weak X-ray emission of these particular stars might be linked to a shift in their peak coronal heating below $\\sim1$ million K and to the formation of very cool coronae. The observed \\ion{N}{5} line strengths, and their ratios to fluxes of cooler chromospheric lines, make that appear extremely unlikely. However, only measurements of lines formed at 0.5--1 million K can answer this question for certain, and the needed observations await the launch of future space missions. Support for this work was provided by NASA through grants GO-6083.01-94 and GO-6446.01-95 from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. \\newpage" + }, + "9701/astro-ph9701219_arXiv.txt": { + "abstract": "The beaming pattern of thermal annihilation radiation is broader than the beaming pattern produced by isotropic nonthermal electrons and positrons in the jets of radio-emitting active galactic nuclei which Compton scatter photons from an external isotropic radiation field. Thus blueshifted thermal annihilation radiation can provide the dominant contribution to the high-energy radiation spectrum at observing angles $\\theta \\gtrsim 1/\\Gamma$, where $\\Gamma$ is the bulk Lorentz factor of the outflowing plasma. This effect may account for the spectral features of MeV blazars discovered with the Compton Telescope on the {\\it Compton Gamma Ray Observatory}. Coordinated gamma-ray observations of annihilation line radiation to infer Doppler factors and VLBI radio observations to measure transverse angular speeds of outflowing plasma blobs can be used to determine the Hubble constant. ", + "introduction": "Observations made with the Compton Telescope (Comptel) on the {\\it Compton Gamma-Ray Observatory} ({\\it CGRO}) reveal the existence of a class of gamma-ray blazars displaying spectral energy distributions peaking in a rather narrow energy band centered at a few MeV (Bloemen et al. \\markcite{Bloem95}1995; Blom et al. \\markcite{Blom95}1995). By combining data from 19 September - 3 October 1991 and 27 December 1991 - 10 January 1992, Bloemen et al. (\\markcite{Bloem95}1995) discovered the first ``MeV blazar\" GRO J0516-609, which shows a 3-4$\\sigma$ peak in its $\\nu F_\\nu$ spectrum in the 1 - 10 MeV energy interval. Only upper limits are measured with Comptel in the 0.75-1 MeV and 10-30 MeV range, although contemporaneous observations with the Energetic Gamma Ray Experiment Telescope (EGRET) on {\\it CGRO} indicate the presence of a weak $3\\sigma$ excess at photon energies $E>100$ MeV. The flat spectrum quasars PKS 0506-612 and PKS 0522-611 are both within the 95\\% Comptel contours of GRO J0516-609, but the former source is the preferred identification due to the more precise EGRET localization, which places only PKS 0506-612 within its 95\\% confidence contour. This source is reported at the 3.9$\\sigma$ significance level in the Second EGRET catalog (Thomson et al. \\markcite{Thomp95}1995), but is too weak to be listed as a catalogued EGRET source. The blazar PKS 0208-512 was detected in analysis of Comptel data by Blom et al. (\\markcite{Blom95}1995). The strongest signal from PKS 0208-512 was obtained by combining data from 8-13 May 1993 and 3-14 June 1993, yielding a flux in the 1-3 MeV band of $4.1(\\pm 0.7)\\times 10^{-4}$ photon cm$^{-2}$ s$^{-1}$ and upper limits at lower and higher energies in the Comptel energy range. Contemporary observations with EGRET yield a strong $> 100$ MeV detection of the source 2EG J0210-5051, with PKS 0208-512 located within the 95\\% confidence contour of this source. The peak in the $\\nu F_\\nu$ spectrum of PKS 0208-512 occurs at MeV energies, supporting its classification as an MeV blazar. Williams et al. (\\markcite{Willi95}1995) also report the detection of the gamma-ray source GRO J1753+57 with a spectrum similar to that of the other two MeV blazars. Lack of association of this object with a flat radio spectrum quasar makes this object's classification uncertain, but one should keep in mind that blazars are highly variable at all wavelengths, so that contemporaneous multiwavelength campaigns could be necessary for proper source identification. The peaked emission from MeV blazars has been interpreted as Doppler boosted $e^+e^-$ annihilation radiation produced in collimated jets of plasma ejected from a supermassive black hole (Roland \\& Hermsen \\markcite{Rolan95}1995; Henri, Pelletier, \\& Roland \\markcite{Henri93}1993). The possible importance of $e^+e^-$ processes and $\\gamma$-$\\gamma$ opacity effects has been considered by Blandford \\& Levinson (\\markcite{Bland94}1994), B\\\"ottcher, Mause, \\& Schlickeiser (\\markcite{B\\\"ottc97}1997), and Marcowith, Henri, \\& Pelletier (\\markcite{Marco95}1995). No attempt has been made, however, to account for the significant spectral differences between MeV blazars and the $> 50$ $\\gamma$-ray blazars detected by EGRET (e.g., von Montigny et al. \\markcite{Monti95}1995), which show strong high-energy radiation extending to GeV and, in a few cases (e.g., Macomb et al. \\markcite{Macom96}1996; Schubnell et al. \\markcite{Schub96}1996), TeV energies. In this paper we show that the beaming pattern of thermal Doppler-boosted $e^+e^-$ annihilation radiation is much broader than the beaming pattern produced by nonthermal jet electrons which Compton scatter photons produced outside the jet, whose beaming pattern was recently derived (Dermer \\markcite{Dermer95}1995; Dermer, Sturner, \\& Schlickeiser \\markcite{Derme97}1997). We call this latter process external Compton scattering (ECS) and distinguish it from synchrotron self-Compton (SSC) scattering, which has also been proposed as the radiation process yielding blazar gamma radiation (see Marscher \\& Travis \\markcite{Marsc96}1996 and references therein). If the high-energy emission in blazars is primarily produced by ECS and, moreover, the annihilation luminosity from thermal $e^+e^-$ pairs is significant, then a class of objects with properties similar to the MeV blazars is a natural consequence of orientation effects when viewing at moderate angles with respect to the axis of the outflowing jet. In addition, we show in Appendix A that if the peaks of MeV blazars are indeed due to Doppler-boosted thermal annihilation radiation, then gamma-ray observations to measure Doppler factors $\\cal D$ coordinated with VLBI radio observations to measure transverse angular speeds $\\mu$ of outflowing radio blobs in two or more episodes is sufficient to determine the Hubble constant $H_0$. ", + "conclusions": "In this paper, we have considered implications of the suggestion that the features in MeV blazars arise from thermal annihilation radiation in a relativistic jet, as suggested by previous authors (see Section 1). The existence of a class of MeV blazars can be understood by orientation effects if the beaming patterns of the nonthermal radiation is different from that of the annihilation radiation. If the gamma rays are produced by nonthermal jet electrons scattering photons produced external to the jet, then the annihilation radiation has a broader beaming pattern. This could account for the existence of the class of MeV blazars, and would be consistent with unification scenarios for radio-emitting AGNs (see, e.g., Urry \\& Padovani \\markcite{Urry95}1995 for a recent review). Moreover, if this interpretation is correct, then we predict that radio galaxies will exhibit time variable hardenings and features at MeV energies which would correlate with increasing core dominance. Gamma-ray observations of line features and contemporaneous VLBI measurements of transverse angular speeds from relativistic plasma outflows provide a method to determine the Hubble constant. The radio observations furnish the apparent speeds, and the gamma-ray observations give the Doppler and cosmologically shifted energy of the annihilation line. As noted earlier by Roland \\& Hermsen \\markcite{Rolan95}(1995), measurements of $\\cal D$ and the apparent superluminal speed can be used to determine the angle of the jet with respect to the observer and the bulk Lorentz factor $\\Gamma$ of the outflowing plasma. As shown in Appendix A, multiple contemporaneous observations of the shifted line energies and transverse angular speeds provide a method to determine the Hubble constant. This is an extension of the work of Schlickeiser \\& Dermer (\\markcite{Schli95}1995), who proposed a model-dependent method for measuring $\\cal D$ from the broadband spectral energy distribution of blazars. If the MeV features can be conclusively identified with annihilation line signatures, then the procedure proposed here will provide a more definitive method for determining the Hubble constant from coordinated radio and gamma ray observations. Arbitrary spectral lines emitted from the outflowing plasma jets can also be utilized by this method. The identification of spectral line features with a specific outflowing plasma blob is indicated at gamma-ray energies by time variability. Observations (Wehrle \\markcite{Wehrl96}1996) show that the emergence of a radio emitting feature in high-resolution VLBI maps is preceded by a gamma-ray flare. A better method would be through milliarcsecond imaging of spectral lines, which is presently only feasible at radio and optical wavelengths; high resolution optical imaging requires, however, bright sources of apparent magnitudes $m \\lesssim 9$. If it is established that the peaks of MeV blazars are a consequence of electron-positron plasma jets, then such a result would have important implications for processes which power jets in radio-loud AGNs, requiring energization of particles in the compact environment near the supermassive black hole. Additional measurements with Comptel on {\\it CGRO}, and observations with more sensitive soft gamma-ray telescopes, such as the {\\it INTEGRAL} telescope scheduled for launch early in the next millenium, could resolve the question of the composition of jets in AGNs and provide a test of this model." + }, + "9701/astro-ph9701055_arXiv.txt": { + "abstract": "The velocity-dependent `one-scale' model of Martins \\& Shellard is used to study the evolution of a cosmic string network (and the corresponding loop population) in open universes. It is shown that in this case there is no linear scaling regime and that even though curvature still dominates the dynamics, at late times strings become the main component of the universe. We also comment on the possible consequences of these results. ", + "introduction": "\\label{s-int} Despite the strong theoretical prejudices favouring a flat universe\\cite{kt,li}, there is a fair amount of observational data which suggests the possibility of an open universe, with a present density that could be as low as $\\Omega\\sim0.3$---notably, the so-called `age problem'\\cite{age} and recent measurements of the baryonic content of x-ray clusters\\cite{xray}. Since in a low-$\\Omega$ universe structures collapse earlier, observations of galaxies at high redshift would also be easier to explain if the universe is open. It is therefore appropriate to consider how some of the standard cosmological scenarios would change if this possibility turns out to be true. One relevant case is that of the evolution of a network of cosmic strings\\cite{vs}. It has been shown\\cite{def,def1} that defect models normalized to COBE in an open universe predict a galaxy power spectrum consistent with that inferred from galaxy surveys without requiring an extreme bias (in general, $\\Omega=1$ models predict more small-scale power than low-$\\Omega h$ ones). However, these results were established either using {\\em a priori} scaling assumptions for the string network\\cite{def} or numerical simulations of texture evolution\\cite{def1}. Here we study the evolution of a cosmic string network in open universes, using the velocity-dependent `one-scale' model of Martins \\& Shellard\\cite{ms,ms1}, which provides the first quantitative description of the complete evolution of the large-scale properties of a cosmic string network. This is briefly summarised below, and used in the following section to obtain the evolutionary properties of both the long-string and the loop populations in an open universe; these are then compared with the standard flat universe case, and some implications of these results are discussed. ", + "conclusions": "\\label{s-cc} In this paper we presented the first discussion of cosmic string evolution in open universes, in the context of the generalized `one-scale' model of Martins \\& Shellard\\cite{ms1}. We have shown that there is no linear scaling regime in an open universe, and that although the string density always decreases with respect to the critical density, it has been increasing relative to that of the background from $z\\sim\\Omega_o^{-1}$, and it will become the main component of the universe sometime in the future. These differences with respect to the standard (flat universe) case only become significant fairly late in the matter-dominated epoch, so with respect to the string-seeded structure formation scenario we should only expect changes on very large scales---that one can easily estimate to be larger than the scales of the largest existing surveys. On the other hand, such large scales are of course relevant when one is comparing the cosmic microwave background anisotropies produced by cosmic strings to COBE data---thereby normalizing the string mass per unit length\\cite{acssv}---since this essentially involves an integration from the present time to the surface of last scattering. Thus the changes in the string network properties discussed in the present paper can significantly alter this normalization. This issue will be discussed in a forthcoming publication." + }, + "9701/astro-ph9701113_arXiv.txt": { + "abstract": "We describe PTreeSPH, a gravity treecode combined with an SPH hydrodynamics code designed for massively parallel supercomputers having distributed memory. Our computational algorithm is based on the popular TreeSPH code of Hernquist \\& Katz (1989). PTreeSPH utilizes a domain decomposition procedure and a synchronous hypercube communication paradigm to build self-contained subvolumes of the simulation on each processor at every timestep. Computations then proceed in a manner analogous to a serial code. We use the Message Passing Interface (MPI) communications package, making our code easily portable to a variety of parallel systems. PTreeSPH uses individual smoothing lengths and timesteps, with a communication algorithm designed to minimize exchange of information while still providing all information required to accurately perform SPH computations. We have additionally incorporated cosmology, periodic boundary conditions with forces calculated using a quadrupole Ewald summation method, and radiative cooling and heating from a parameterized ionizing background following Katz, Weinberg \\& Hernquist (1996). The addition of other physical processes, such as star formation, is straightforward. A cosmological simulation from $z=49$ to $z=2$ with $64^3$ gas particles and $64^3$ dark matter particles requires $\\sim 6000$ node-hours on a Cray T3D, with a communications overhead of $\\sim 10$\\% and is load balanced to a $\\sim 90$\\% level. When used on the new Cray T3E, this code will be capable of performing cosmological hydrodynamical simulations down to $z=0$ with $\\sim 2\\times 10^6$ particles, or to $z=2$ with $\\sim 10^7$ particles, in a reasonable amount of time. Even larger simulations will be practical in situations where the matter is not highly clustered or when periodic boundaries are not required. \\bigskip \\bigskip \\noindent {\\it Subject Headings:} methods: numerical --- cosmology: theory \\newpage ", + "introduction": "Gas dynamics and gravitational forces together govern the evolution of astrophysical systems on nearly all scales. Star and planet formation are thought to occur in accretion disks well-represented as collapsing gas clouds in quasi-hydrostatic equilibrium. The formation and evolution of star clusters is governed by interactions with the interstellar medium, as well as viscous forces generated from tidal perturbations in individual stellar interactions. Shock heating from supernovae feedback is important for the observed structure of galaxies and the creation of hot X-ray halos around clusters. All the galaxies we see are themselves thought to have resulted from dissipational collapse through heirarchical gravitational instability. On these largest scales, comparison between observation and theory requires some understanding of the interplay between dissipationless dark matter and baryonic material. Analytical treatments of the processes relevant to cosmology are generally applicable only when the matter is still evolving in a linear or quasi-linear fashion, or more generally for objects possessing inherent symmetries such as stars, supernovae, or accretion disks. However, many systems are intrinsically asymmetric and/or highly nonlinear, such as galaxies forming from primordial fluctuations. Semi-analytical treatments of galaxy formation invariably make {\\it ad hoc} assumptions about the relationships between gas and dark matter in the highly nonlinear regime (\\eg \\cite{kau93}; \\cite{hey95}; \\cite{som97}). A proper understanding of galaxy and structure formation requires that dissipation, shocks, and pressure forces be taken into account, in addition to gravitational collapse. Highly nonlinear systems are best modelled numerically by evolving large numbers of particles and/or grid cells self-consistently under both gravitational and hydrodynamical forces. A variety of techniques to do this have been developed, each with its own advantages and disadvantages. Gravitational forces are computed using either grid-based or particle-based algorithms. The simplest particle-based technique directly sums pairwise forces between all particle pairs. These codes are useful for many purposes (\\eg \\cite{aar85}), but the computational time grows with particle number $N$ as $O(N^2)$ or worse, making them appropriate only for small systems ($N\\la 10^4$). Tree codes place particles in a hierarchical data structure, and use multipole expansions to approximate the force between distant groups of particles, reducing the scaling to $O(N\\log{N})$ (\\eg Barnes \\& Hut 1986, hereafter \\cite{bar86}). In particle-mesh (PM) codes, the gravitational potential is computed on a grid using Fast Fourier Transforms (FFTs), which also scales as $O(N\\log{N})$ but with fewer operations than a treecode (\\eg \\cite{hoc80}). The main drawback of this approach is that the resolution is limited by the cell size. Adaptive mesh codes have been developed to subdivide or deform cells in dense regions to obtain better dynamic range (\\cite{vil89}; \\cite{pen95}; \\cite{xu97}). Also, hybrid codes such as PP-PM (P$^3$M; \\cite{efs85}) and Tree-PM (TPM; \\cite{xu95}) alleviate resolution limitations. Other techniques, such as the self-consistent field (SCF) method (\\eg \\cite{her92}) can be even faster, but are generally useful only for systems not far from a well-specified equilibrium. Hydrodynamical forces can be computed in either a Lagrangian or Eulerian manner. Eulerian codes represent the fluid on a grid of cells, and compute the flux of fluid across cell boundaries, as in the Piecewise Parabolic Method (PPM) (\\cite{col84}; \\cite{bry94}). In Lagrangian codes, the dynamical equations are obtained from the Lagrangian form of the hydrodynamical conservation laws. Some Lagrangian codes represent the fluid by particles without the use of a grid, as in Smoothed Particle Hydrodynamics (SPH) (\\cite{luc77}; \\cite{gin77}). In principle, any gravity solver may be combined with any hydrodynamics method. Grid-based codes include PM-PPM (\\cite{bry94}), PM-TVD (\\cite{ryu93}), and adaptive mesh hydrodynamics codes (\\eg \\cite{bry95}). Lagrangian examples include P$^3$M-SPH (\\cite{evr88}; \\cite{cou95}), and GRAPE-SPH (\\cite{ste96}), which uses the special-purpose GRAPE hardware to perform rapid pairwise gravity summation. Lagrangian codes provide much better spatial resolution in high density regions compared with Eulerian codes for a given computational expense, at the cost of poorer shock resolution and lower resolution in underdense regions (\\cite{kan94}). For many astrophysical applications the overdense regions are of most immediate interest, and in those cases Lagrangian codes are preferable if shocks do not dominate the dynamics of the system. In this paper we focus on a pure particle-based combination of gravity and hydrodynamics solvers analogous to the TreeSPH code of Hernquist \\& Katz (1989; hereafter \\cite{her89}). In SPH, the gas is sampled and represented by particles, which are smoothed to obtain a continuous distribution of gas properties. Since there is no grid, there are no inherent constraints on the global geometry or spatial resolution. Neighbor finding is also done using a tree structure, and thus the entire code scales as $\\sim O(N\\log N)$. However, unlike grid-based codes, SPH cannot handle arbitrarily large gradients due to its finite particle resolution. Also, an artificial viscosity is used to capture shocks, further limiting the spatial resolution locally. Despite these compromises, TreeSPH has been successfully used in a wide range of astrophysical applications, including giant molecular clouds (\\cite{gam92}), colliding galaxies (\\cite{mih94}), ram pressure stripping in clusters (\\cite{kun93}), formation of galaxies (\\cite{kat91}) and cluster (\\cite{kat93}), and the high-redshift Lyman alpha forest (\\cite{her96}). A major goal in numerical astrophysics is to improve the dynamic range of simulations. One would ideally like to simulate volumes comparable to the size of the Universe ($\\sim 10^3$ Mpc), but resolve star forming regions in galaxies on the parsec scale. The spatial dynamical range required per dimension is thus $\\sim 10^9$, well beyond the $\\sim 10^3$ in dynamic range that codes can currently achieve. High resolution studies of even single collapsing protogalaxies in a cosmological setting require a dynamic range $\\ga 10^4$. TreeSPH endows each gas particle with its own smoothing length and timestep, thus improving the dynamic range substantially by being adaptive in both space and time. TreeSPH has been vectorized efficiently (\\cite{her89}; \\cite{her90}), but even on vector supercomputers the largest TreeSPH cosmological simulation to date (Katz, Weinberg \\& Hernquist 1996, hereafter \\cite{kat96}) employed a total of $\\approx 524,288$ particles with a dynamic range of $\\sim 2000$. Massively parallel supercomputers (MPPs) link hundreds of workstation processors together to yield an overall computational power more than an order of magnitude greater than that of current vector supercomputers. While MPPs are attractive, there are a number of major difficulties in adapting codes to run on these machines, often requiring significant algorithmic changes from serial or vector code. The first difficulty is that in distributed memory systems, each processor possesses only a relatively small amount of local memory, and accessing information from another processor's memory is slow compared to the computation speed. Thus a parallel code must subdivide a simulation and exchange information between processors in a manner which minimizes communication time, while not taxing each processor's memory. The second complication is {\\it load balancing}, \\ie insuring that no processor spends a significant amount of time idly waiting for another processor to send required information. Good load balancing can be achieved either by designing an algorithm so that each processor has roughly an equal amount of computational work between {\\it synchronous} communications, or by implementing an {\\it asynchronous} communication scheme by which processors continue to do other computations while waiting to send or receive information. In this paper we present a parallel implementation of a gravitational treecode combined with SPH, called PTreeSPH. While the underlying numerical techniques are similar to those in TreeSPH, our implementation on MPP machines required a complete redesign of the code as well as several major algorithmic changes. PTreeSPH is a $C$ code which uses a domain decomposition prescription to subdivide the simulation and a synchronous hypercube message-passing paradigm to build small ``locally essential\" simulation subvolumes on each processor. The $N$-body portion of the code was developed by Dubinski (1996; hereafter \\cite{dub96}), to which we have added a parallelized version of SPH as well as the dynamics required for cosmological simulations. Building locally essential simulations on each processor allows the parallelization to be decoupled from the computations, making it straightforward to incorporate additional physical processes. In addition, information transfers occur in single bursts rather than continually during a simulation, thereby lowering communication overhead. We utilize the Message Passing Interface (MPI) package exclusively to handle processor communications, which makes the code easily portable to most major parallel supercomputers, including the Cray T3D/T3E, IBM SP2, and Intel Paragon, shared memory systems such as the SGI Power Challenge, as well as to networks of workstations. While our parallel algorithm could likely be made more efficient by tailoring it to specific machines and including some asynchronous communication, we are more interested in producing a reasonably efficient code which is portable and adaptable to a wide variety of applications. Several other groups are now parallelizing SPH in various forms. The Virgo Consortium has recently developed a parallel P$^3$M-SPH code capable of doing cosmological simulations (\\cite{pea95}; \\cite{jen96}). Evrard and Brieu (private communication) are working on a similar code. \\cite{war94} have developed a generalized parallel $N$-body code using a hashed octal tree structure to asynchronously access information on other processors, and they describe how their code may be applied to neighbor finding as required in SPH. Also, \\cite{dik97} have written a parallel $N$-body treecode similar in spirit to \\cite{war94}, and are now adding their own parallel version of SPH to it. ", + "conclusions": "A parallel TreeSPH code for performing cosmological hydrodynamical simulations has been presented. The physics is incorporated in a manner quite similar to the TreeSPH code (\\cite{her89}; \\cite{kat96}), but the algorithm has been redesigned to handle difficulties unique to the massively parallel environment. The code is capable of performing cosmological simulations in a periodic volume, including the effects of radiative cooling and heating from a parameterized ionizing background. The use of massively parallel supercomputers currently enables one to perform simulations with roughly an order of magnitude more particles than is practically possible on a vector supercomputer. Moreover, with the current rate of advance in chip technology, the massively parallel approach promises to provide rapid further increases in computational power and therefore simulation size. Currently, a cosmological hydrodynamical simulation of $10^7$ particles down to $z=2$ is feasible within a typical allocation of supercomputer time; alternatively, one can explore models with many smaller simulations. Significantly larger simulations are possible for applications which do not require periodic boundaries, or whose matter is less clustered than in a cosmological simulation. PTreeSPH has been implemented with primary concern for portability and expandability, and secondary concern for optimal efficiency. Nevertheless, the code achieves good load balancing ($\\sim 90$\\%) and has fairly low communication overhead ($\\sim 10$\\%). Compile-time flags have been implemented to easily include or exclude various physical processes such as hydrodynamics, cooling, and cosmology. The use of the MPI message passing software allows for easy portability between parallel platforms; already the code has been run successfully on the Cray T3D, the IBM SP2, and an SGI Power Challenge. The building of locally essential problems, both for gravity and hydrodynamic forces, allows for straightforward incorporation of other physical processes as desired. Together with its fully Lagrangian nature, this will make PTreeSPH useful for a wide variety of astrophysical applications, including cosmology, galaxy interactions and cluster formation." + }, + "9701/astro-ph9701107_arXiv.txt": { + "abstract": "The {\\RO} WFC survey has provided us with evidence for the existence of a previously unidentified sample of hot white dwarfs (WD) in non-interacting binary systems, through the detection of EUV and soft X-ray emission. These stars are hidden at optical wavelengths due to their close proximity to much more luminous main sequence (MS) companions (spectral type K or earlier). However, for companions of spectral type $\\sim$A5 or later the white dwarfs are easily visible at far-UV wavelengths, and can be identified in spectra taken by {\\IUE}. Eleven white dwarf binary systems have previously been found in this way from {\\RO}, {\\euve} and {\\IUE} observations (e.g. Barstow et al. 1994). In this paper we report the discovery of three more such systems through our programmes in recent episodes of {\\IUE}. The new binaries are HD2133, RE J0357$+$283 (whose existence was predicted by Jeffries, Burleigh and Robb 1996), and BD$+$27$^\\circ$1888. In addition, we have independently identified a fourth new WD$+$MS binary, RE J1027$+$322, which has also been reported in the literature by Genova et al. (1995), bringing the total number of such systems discovered as a result of the EUV surveys to fifteen. We also discuss here six stars which were observed as part of the programme, but where no white dwarf companion was found. Four of these are coronally active. Finally, we present an analysis of the WD$+$K0IV binary HD18131 (Vennes et al. 1995), which includes the {\\RO} PSPC X-ray data. ", + "introduction": " ", + "conclusions": "" + }, + "9701/astro-ph9701088_arXiv.txt": { + "abstract": "We test through stellar N-body simulations some scenarios to explain the dynamics of the peculiar nucleus of the Andromeda galaxy (M~31): although HST observations reveal a double nucleus morphology, the rotation field is almost symmetric around the bulge gravity centre and the velocity dispersion is off-centred. We show that any $m=1$ perturbation has a very short life-time (a few 10$^5$ yr). Assuming that the bright peak (P1) is a cold stellar cluster infalling into the nucleus, and that the large central velocity gradient is due to a central dark mass (in the range 7~$10^7$--$10^8$~\\Msun), we obtain a reasonably good fit to the observations. However, if this cluster lies in the central 20 pc, we estimate the life-time of the cluster to be less than 0.5~Myr. The dynamical friction is more efficient than estimated by analytic formulae, and is essentially due to the deformation of the stellar cluster through the huge tidal forces provided by the black hole. We show that the cluster cannot be on a circular orbit around the centre if the nucleus hosts a massive black hole of a few $10^7$~\\Msun, and finally provide some estimates of the kinematics as observed with HST. ", + "introduction": "Each new observational result concerning the nucleus of M~31 seems to deepen the mystery of its structure. This is mainly due to the improvement in the achieved spatial resolution of the photometric and spectroscopic data: the better resolved the object is, the more complex it appears. For any proposed physical mechanism which can quantitatively reproduce the observables, it is important to provide some specific predictions. In this respect, the nucleus of M~31 is certainly an excellent laboratory to test our knowledge of galactic nuclei, as we can now optically resolve scales as small as 0.3 pc (e.g. HST or optically adaptive ground-based systems). Indeed, the numerous studies achieved on this object revealed a puzzling complexity. M~31's nucleus is a compact stellar system with an average ellipticity of $\\ave{\\epsilon} \\sim 0.4$. Asymmetries in its major-axis surface brightness profile (in the visible) were already observed in 1974 by Light et al., and subsequently understood using the pre COSTAR HST data as a double component nucleus (Lauer et al. 1993). In the V band, the fainter peak (P2) is almost coincident with the centre of the bulge isophotes, as the brightest one (P1) is located at about $0\\farcs5$ ($\\sim 1.8$ pc for a distance to M~31 of 0.77 Mpc) from P2 roughly along the major-axis. At 175 nm however, P2 is the brightest point, with P1 having an UV upturn similar to its surroundings (King et al. 1995). In 1960, Lallemand et al. obtained spectrograms of the central region of M~31 and detected the rapid rotation of its nucleus, concluding that it is a dynamically independent structure. This was confirmed by Kormendy (1988) and Dressler \\& Richstone (1988) by long-slit spectroscopy with CCD detectors. Both studies suggest the presence of a central dark mass of the order of $10^7$ M$_{\\odot}$ in the centre of the nucleus which would account for the high value of the central velocity and velocity dispersion gradients. Sub-arcsecond velocity and dispersion maps were obtained with the TIGER spectrograph by Bacon et al. (1994, hereafter BEMN94) which uncovered more asymmetries (already traced in previous published data): the velocity field is nearly symmetric about a point $V_0$ located very close to P2 ($0\\farcs05$ according to Lauer et al. 1993), but the maximum $S$ of the stellar velocity dispersion roughly corresponds to the symmetric point of P1 with respect to P2 (see Fig.~26 of BEMN94). Different interpretations have been proposed to explain the observed photometric and dynamical asymmetries in the nucleus of M~31 (see Sect.~\\ref{sec:meca} and references therein), but they each encounter severe problems. Some questions remain open, such as the actual efficiency of dynamical friction for an external body falling into the nucleus, the life-time of $m=1$ perturbations, or the amount of dynamical perturbation on the nucleus itself; we try here to answer these questions through N body simulations. In this paper, we concentrate on one of the proposed mechanism, namely that P1 corresponds to a stellar cluster falling into the potential well of the nucleus, hosting a massive black hole. In a companion paper, we will examine whether a central black hole is unavoidable, or the presence of a nuclear bar could account for the observational data. We summarize and comment the alternatives already discussed by different authors in Sect.~2. In Sect.~3 we describe the N body code we used as well as the initial conditions of the different experiments. The results are discussed in Sect.~4. In Sect.~5, we present and discuss our attempts to fit the observables using different assumptions. The general discussion and conclusions are given in Sect.~6. ", + "conclusions": "\\label{sec:disc} In this paper we have presented N-body simulations of a stellar cluster falling towards a galactic nucleus. These experiments were designed to ressemble the central region of M~31. First we have shown that it was possible to obtain reasonable fits to all observables, either with a projected stellar cluster (which served to adapt our initial conditions for the N body simulations), or with a stellar cluster close to the nucleus and interacting with it. The best model considered an initial orbit where the cluster was unbound to the nucleus. Slight discrepancies between this model and the data still remain (e.g. the nucleus isophotes flattening). This solution is certainly not unique and other significantly different initial conditions (e.g. the choice of the orbit) could lead to very similar observables. Second, we have shown that the dynamical friction is even more efficient than previously estimated: the cluster rapidly decays to smaller distances, where it is disrupted through tidal interaction with the BH. The lifetime of such a cluster is short ($< 5 \\; 10^5$) at radii smaller than $\\sim 10$~pc if a central dark mass of a few $10^7$~\\Msun is present at the centre. The main argument against the scenario of a falling cluster is that we must see the M31 nucleus at a very special time. However, the decaying of a stellar cluster is not a rare phenomenon. Tremaine et al (1975) proposed that the falling of globular clusters could even be the formation mechanism of the nuclear disc. Orders of magnitude are consistent with the observations, if about 25 clusters have been disrupted to form the nucleus, which mass has grown as $t^{1/2}$. This means that one cluster is falling every 400 Myr; since it will perturb the brightness distribution during $\\approx$ 0.2 Myr, the probability of the present configuration of the M31 nucleus is about 4 10$^{-4}$. This however includes only globular clusters spiraling in on a quasi-circular orbit. We have shown that P1 being on a circular orbit is inconsistent with the presence of a central dark mass of a few $10^7$~\\Msun. More rapid globular clusters, in hyperbolic orbits, as in one of the scenarii proposed here, would be disrupted, but not slowed down enough to contribute to the nucleus mass. Therefore, we should conclude that the hypothesis of a falling cluster (e.g. nucleus of an accreted dwarf galaxy, globular cluster) remains a viable explanation. Lauer et al. (1993) suggested that P1 could survive close to the centre if it contains a secondary black hole of ``substantial mass''. The probability of seeing it now in this configuration is even smaller than for a stellar cluster, given the frequencies of black holes. Only if we assume that dark matter in galactic halos is constituted by BHs are the latter more numerous than globular clusters. In the model of Lacey \\& Ostriker (1985), one of these 10$^6$ M$_\\odot$ BH arrives in the centre every 10$^8$ yrs, and most of the time it merges with the BH already present, or is ejected if there is already a binary BH. Xu \\& Ostriker (1994) argue that, since fresh BH arrive regularly on the centre, the probability to find a binary BH with a separation of 1.6pc (as P1-P2) is 10\\%. It is however difficult to explain why P1 is so bright, without any cusp contrary to P2. More stringent is the constraint of a low dispersion at P1: in the case of a secondary black hole at P1, the dispersion there should be significantly larger than observed. Our simulations finally served to precise the required BH mass, in the hypothesis of an axisymmetric nucleus: in order to fit the disc thickness and stability requirements, we need the presence of a mass concentration of at least $7\\;10^7$~\\Msun to explain the central stellar velocity and velocity dispersion gradients. The detection of a strong central UV and X ray source in the centre of M~31 is also in favor of a supermassive black hole. The main underlying assumption in the kinematical determination of the central dark mass is the {\\em axisymmetry of the nucleus}. In a forthcoming paper, we will examine the possibility of a triaxial morphology (i.e. a bar) for the nucleus of M~31." + }, + "9701/astro-ph9701041_arXiv.txt": { + "abstract": "We present a new ultra-violet spectrum of the QSO 0013$-$004 with 0.9 \\AA\\ resolution obtained with the MMT Blue spectrograph. The $\\upsilon$ = 0 - 0, 1 - 0, 2 - 0 and 3 - 0 Lyman bands of H$_2$ associated with the z = 1.9731 damped Ly\\al\\ absorption line system have been detected. The H$_2$ column density is N(H$_2$) = 6.9$(\\pm 1.6)\\times 10^{19}$ \\cm, and the Doppler parameter b = 15$\\pm 2$ \\kms. The populations of different rotational levels are measured and used to derive the excitation temperatures. The estimated kinetic temperature T$_K\\sim 70$ K, and the total particle number density n$(H) \\sim$ 300 cm$^{-3}$. The UV photoabsorption rate $\\beta_0 \\sim 6.7\\times 10^{-9}$ s$^{-1}$, about a factor of few times greater than that in a typical diffuse Milky Way interstellar cloud. The total hydrogen column density is N(H) = 6.4($\\pm 0.5)\\times 10^{20}$ \\cm. The fractional H$_2$ abundance $f$ = 2N(H$_2$)/(2N(H$_2$) + N(H I)) $\\sim$ 0.22 $\\pm$ 0.05 is the highest among all observed damped Ly\\al\\ absorbers. The high fractional H$_2$ abundance is consistent with the inferred presence of dust and strong C I absorption in this absorber. ", + "introduction": "Damped Ly$\\alpha$ quasar absorption line systems are commonly believed to trace a population of objects which may be the progenitors of modern galaxies (e.g. Wolfe 1990). The typical metallicity and dust-to-gas ratio in the damped Ly$\\alpha$ systems at z $\\sim$ 2 is measured to be about 10\\% of those in the Milky Way with a large scatter (e.g. Pettini {\\em et al.} 1994; 1995; Pei, Fall \\& Bechtold, 1991; Lu {\\em et al.} 1996). The UV Lyman and Werner bands of H$_2$ have been difficult to detect however, so that little is known about molecular gas in damped Ly\\al\\ galaxies (Levshakov {\\em et al.} 1992; Lanzetta {\\em et al.} 1989; Foltz et al 1988; Chaffee {\\em et al.} 1988; Black {\\em et al.} 1987). The problem such searches face is that without high resolution and high signal-to-noise ratio spectra, the molecular hydrogen lines are blended with the Ly\\al\\ forest lines. Further, quasar absorption line studies usually concentrate on bright blue quasars, so lines of sight with substantial reddening and large molecular fraction are probably selected against (Ostriker \\& Heisler, 1984; Fall \\& Pei 1993). So far, there is only one detection of H$_2$ absorption in the z = 2.811 damped Ly\\al\\ absorber toward PSK 0528$-$250, which shows moderate H$_2$ absorption with total column density N(H$_2$) = $1\\times 10^{18}$ \\cm\\ (Foltz {\\em et al.} 1988; Lanzetta 1993). This system may not be representative since its redshift is very close to the redshift of the quasar. Previous observations of the z$_{ab}$ = 1.9731 damped Ly\\al\\ absorber toward Q 0013$-$004 (z$_{em}$ = 2.0835) showed C I absorption, and line ratios of Cr II, Fe II and Zn II indicated significant depletion, and hence a relatively high dust-to-gas ratio compared to other damped Ly\\al\\ absorbers (Pettini et al. 1994, Ge, Bechtold, \\& Black 1996). Thus, it seemed to be a good system to search for H$_2$. In this {\\it Letter}, we report the detection of strong H$_2$ absorption lines. ", + "conclusions": "The molecular fraction $f$ = 0.22 $\\pm$ 0.05 for the z = 1.97 absorber is similar to that seen in E(B $-$ V) $>$ 0.1 clouds in the Milky Way (Savage {\\em et al.} 1977). Indeed, the upper limit of the relative depletion of Cr to Zn, [Cr/Zn] $\\le$ $-$ 1.0 (Pettini {\\em et al.} 1994) suggests that the dust-to-gas ratio in this absorber is much higher than the average value of 10\\% of the Milky Way's ratio for damped systems at z $\\sim$ 2 (Pettini {\\em et al.} 1994). The previous observation of Fe II \\lam\\ 1608.45 \\AA\\ provides a direct comparison of the relative depletion of Fe to Zn, [Fe/Zn] = $-$ 1.2 (Ge {\\em et al.} 1997) to those of the Milky Way's diffuse clouds. Table 3 shows a comparison between the z = 1.9731 absorber and the diffuse clouds toward $\\zeta$ Oph and $\\xi$ Per (Spitzer {\\em et al}. 1974; Savage {\\em et al.} 1977; Jura 1975; Savage {\\em et al} 1992; Cardelli {\\em et al.} 1991). The relative depletion of [Fe/Zn] implies that the dust-to-gas ratio in the z = 1.9731 absorber is $\\sim$ 50\\% of the Milky Way's value. The relatively normal dust-to-gas ratio is consistent with the high H$_2$ fraction because H$_2$ is very efficiently formed on dust grain surfaces (Savage {\\em et al.} 1977). As discussed by previous reviews (e.g. Spitzer \\& Jenkins, 1975; Shull \\& Beckwith 1982), the relative populations of J = 0 and 1 are established dominantly by thermal particle collisions, especially for saturated lines, so the excitation temperature, T$_{01}$, is approximately equal to the kinetic temperature, T$_k$, of the clouds. The measured T$_{01} = 70\\pm 13$ K in the z = 1.9731 absorber implies that the kinetic temperature is similar to that of the Milky Way diffuse clouds, $\\langle T_{01}\\rangle$ = (77 $\\pm$ 17) K (Savage {\\em et al.} 1977). The higher rotational levels are populated primarily by collisions, formation pumping and UV pumping and radiative cascade after photoabsorption to the Lyman and Werner bands (e.g. Spitzer \\& Zweibel 1974; Jura 1974, 1975). For example, the J = 4 is populated by direct formation pumping, and by UV pumping from J = 0 (e.g. Black \\& Dalgarno 1976). For densities less than $10^4$ cm$^{-3}$, the J = 4 level is depopulated mainly by spontaneous emission (Dalgarno \\& Wright 1972; Elitzur \\& Watson 1978). Therefore, in a steady state for J = 4, \\begin{equation} p_{4,0}\\beta(0)n(H_2,J = 0) + 0.19 Rn(H)n = A_{42}n(H_2, J = 4), \\end{equation} where $p_{4,0} = 0.26$ is the UV pumping efficiency into the J = 4 level from the J = 0 level (Jura 1975), $A_{42} = 2.8\\times 10^{-9}$ s$^{-1}$ is the spontaneous transition probability (Dalgarno \\& Wright 1972), R is the H$_2$ formation rate, n = n(H) + 2n(H$_2$), and $\\beta(J)$ denotes the rate of absorption in the Lyman and Werner bands from the Jth rotational level, including any attenuation (see Jura 1975 for details). The equilibrium between the H$_2$ formation on dust grains and H$_2$ destruction by absorption of Lyman- and Werner-band radiation (e.g. Jura 1975) can be written as \\begin{equation} I n(H_2) = R n(H)n \\approx 0.11 \\sum_{J=0}^6 \\beta(J) n(H_2,J), \\end{equation} where $I$ is the H$_2$ dissociation rate (Jura 1975). If the self-shielding in J = 0, 1 levels is about the same, so that $\\beta(0)\\approx \\beta(1)$, then \\begin{equation} n(H_2,J=4)A_{42} = 1.52 Rn(H)n. \\end{equation} Thus, $Rn = 8.1(^{+3.9}_{-2.6})\\times 10^{-15}$ s$^{-1}$ for the Q 0013$-$004 cloud, which is about the same magnitude as that for $\\xi$ Per and $\\zeta$ Oph clouds (Jura 1975). We can use the analytic calculation for n(H$_2$)/n(H) within a H$_2$ cloud by Jura (1974) to estimate the H$_2$ dissociation rate $I \\approx 7.4\\times 10^{-10}$ s$^{-1}$. The photoabsorption rate in the Lyman and Werner bands outside of the cloud, $\\beta_0 \\approx I/0.11 \\approx 6.7 \\times 10^{-9}$ s$^{-1}$, which is similar to that of $\\xi$ Per cloud (Jura 1975) and about a factor of a few higher than that of $\\zeta$ Oph cloud (Federman {\\em et al.} 1995). Further, the photoabsorption rate, $\\beta_0$, depends linearly on the local radiation field at 930 - 1150 \\AA\\ (Jura 1974), therefore, the $\\beta_0$ for the z = 1.9731 absorber corresponds to an estimated local radiation field at 1000 \\AA\\ of J$_{1000 \\AA}$ $\\approx 3\\times 10^{-18}$ ergs cm$^{-2}$s$^{-1}$Hz$^{-1}$ster$^{-1}$. Thus J$_{1000\\AA}$ is about three orders of magnitude higher than the radiation field at the Lyman limit, J$_{912\\AA}$, expected in the ambient IGM at this redshift (e.g. Lu {\\em et al.} 1991; Bechtold 1994). The radiation field at 1000 \\AA\\ is therefore probably dominated by the UV emission by hot stars in this galaxy. As mentioned above, the z = 1.9731 absorber has a similar dust-to-gas ratio to that of the Milky Way diffuse clouds. If we assume the H$_2$ formation rate on grains, R $\\sim 3\\times 10^{-17}$ cm$^{3}$s$^{-1}$, the typical rate for the Milky Way's clouds (Jura 1975), then the inferred value of the number density n $\\sim$ 300 cm$^{-3}$, about the same as that of the $\\zeta$ Oph and $\\xi$ Per clouds. The derived values of the UV radiation field, density, and temperature in the z = 1.97 absorber are estimates based on the simple analysis of Jura (1975), which ignores the depth-dependence of the attenuation by dust and self-shielding of absorption lines. However, the results from this simple analysis are qualitatively consistent with that from more detailed modeling (e.g. van Dishoeck \\& Black 1986). We note that there is some uncertainty in the derived b-value. While b = 15 \\kms\\ provides the best fit to the observed values, smaller Doppler parameters, for example, b = 2 \\kms, also give an acceptable solution. In this case, the implied column densities in J = 2, 3 and 4 are so high they would suggest that conditions more like a photon-dominated region (PDR) are implied (e.g. Draine \\& Bertoldi 1996; Black \\& van Dishoeck 1987; Abgrall {\\em et al.} 1992; Le Bourlet {\\em et al.} 1993; Sternberg \\& Dalgarno 1989). The populations in J = 2, 3 and 4 would be fit by a single excitation temperature of 350 K. This would leave excess population in J = 0 and 1, suggesting a cold component at T$_{01}\\approx 63$ K. The density of the absorber cloud would have to be lower than $\\sim$ 5000 cm$^{-3}$ in order for the J = 5 limit to be consistent with the populations in J = 2, 3 and 4. Finally, if the larger value b = 15 \\kms\\ is correct, then this absorber has a b-value which is much larger than that for typical Milky Way diffuse clouds (e.g. Spitzer {\\em et al.} 1974). This would indicate that there is likely more than one velocity component. Spectra of higher resolution would permit a better constrained analysis of the excitation and molecular abundance and provide a better understanding of physical conditions in this high redshift galaxy." + }, + "9701/astro-ph9701097_arXiv.txt": { + "abstract": "A recent homogeneous study of outflow activity in low-mass embedded young stellar objects (YSOs) (Bontemps et al. 1996) suggests that mass ejection {\\it and} mass accretion both decline significantly with time during protostellar evolution. In the present paper, we propose that this rapid decay of accretion/ejection activity is a direct result of the non-singular density profiles characterizing pre-collapse clouds. Submillimeter dust continuum mapping indicates that the radial profiles of pre-stellar cores flatten out near their centers, being much flatter than $\\rho(r) \\propto r^{-2}$ at radii less than a few thousand AU (Ward-Thompson et al. 1994). In some cases, sharp edges are observed at a finite core radius. Here we show, through Lagrangian analytical calculations, that the supersonic gravitational collapse of pre-stellar cloud cores with such centrally peaked, but flattened density profiles leads to a transitory phase of energetic accretion immediately following the formation of the central hydrostatic protostar. Physically, the collapse occurs in various stages. The first stage corresponds to the nearly isothermal, dynamical collapse of the pre-stellar flat inner region, which ends with the formation of a finite-mass stellar nucleus. This phase is essentially non-existent in the `standard' singular model developed by Shu and co-workers. In a second stage, the remaining cloud core material accretes supersonically onto a non-zero point mass. Because of the significant infall velocity field achieved during the first collapse stage, the accretion rate is initially higher than in the Shu model. This enhanced accretion persists as long as the gravitational pull of the initial point mass remains significant. The accretion rate then quickly converges towards the characteristic value $\\sim a^3/G$ (where $a$ is the sound speed), which is also the constant rate found by Shu (1977). If the model pre-stellar core has a finite outer boundary, there is a terminal decline of the accretion rate at late times due to the finite reservoir of mass.\\\\ We suggest that the initial epoch of vigorous accretion predicted by our non-singular model coincides with Class~$0$ protostars, which would explain their unusually powerful jets compared to the more evolved Class~I YSOs. We use a simple two-component power-law model to fit the diagrams of outflow power versus envelope mass observed by Bontemps et al. (1996), and suggest that Taurus and $\\rho$ Ophiuchi YSOs follow different accretion histories because of differing initial conditions. While the isolated Class~I sources of Taurus are relatively well explained by the standard Shu model, most of the Class~I objects of the $\\rho$ Oph cluster may be effectively in their terminal accretion phase. ", + "introduction": "Despite recent observational and theoretical progress, the initial conditions of star formation and the first phases of protostellar collapse remain poorly known. It is reasonably well established that low-mass stars form from the collapse of centrally-condensed cloud cores initially supported against gravity by a combination of thermal, magnetic, and turbulent pressures (e.g. Shu et al. 1987, 1993 for reviews). However, the critical conditions beyond which a cloud core becomes unstable and starts to collapse are uncertain and still a matter of debate (e.g. Shu 1977, Mouschovias 1991, Boss 1995, Whitworth et al. 1996). In particular, they depend on yet unmeasured factors such as the strengths of the static and fluctuating components of the magnetic field.\\\\ Once fast cloud collapse sets in, the main {\\it theoretical} features of the dynamical evolution that follows have been known since the pioneering work of Larson (1969). During a probably brief first phase, the released gravitational energy is freely radiated away and the cloud stays isothermal. This initial collapse phase, which tends to produce a strong central concentration of matter, ends with the formation of an opaque, hydrostatic stellar object (cf. Larson 1969; Boss \\& Yorke 1995). This time is often denoted $t = 0$ and referred to as (stellar) `core formation' in the literature (e.g. Hunter 1977). When the stellar core has fully formed, one enters the {\\it accretion phase} during which the central protostar builds up its mass ($M_{\\star}$) from a surrounding infalling envelope (of mass $M_{env}$) while progressively warming up. The infalling gas is arrested and thermalized in an accretion shock at the surface of the stellar core, generating an infall luminosity L$_{inf} \\approx GM_\\star(t)\\dot{M}_{acc}/R_\\star$. In the `standard' theory of Shu et al. (1993) which uses singular isothermal spheres as initial conditions, the accretion rate $\\dot{M}_{acc}$ is constant and equal to $a_{eff}^3/G$, where $a_{eff}$ is the effective sound speed. The singular isothermal sphere, which has $\\rho \\propto (a_{eff}^2/G)\\, r^{-2}$, is a physically meaningful starting point for the `self-initiated' collapse of an isolated cloud core because it represents the (unstable) limit of infinite central concentration in equilibrium models for self-gravitating isothermal spheres (e.g. Shu 1977 -- hereafter Shu77). Furthermore, Lizano \\& Shu (1989) have shown that magnetically-supported cloud cores undergoing ambipolar diffusion evolve naturally toward a singular configuration reminiscent of a singular isothermal sphere (see also Ciolek \\& Mouschovias 1994). However, it is likely that actual cloud cores become unstable and start to collapse before reaching the asymptotic singular state, especially when they are perturbed by an external agent such as a shock wave (e.g. Boss 1995, Whitworth et al. 1996). In this case, the initial conditions will be effectively non-singular, and $\\dot{M}_{acc}$ is expected to be time-dependent (e.g. Zinnecker \\& Tscharnuter 1984, Henriksen 1994, Foster \\& Chevalier 1993, McLaughlin \\& Pudritz 1997). It is this possibility that we explore further and compare with relevant observations in the present paper.\\\\ An unsettled related issue concerns the manner in which the central hydrostatic stellar core develops. While in some models core formation occurs in a dynamical, {\\it supersonic} fashion (e.g. Larson 1969, Foster \\& Chevalier 1993, present paper), this formation is achieved by slow, {\\it subsonic} evolution of the gas in the scenario advocated by Shu and co-workers (see also Ciolek \\& Mouschovias 1994). An important observational consequence is that while the dynamical models predict the existence of `isothermal protostars' (in the sense of the initial collapse phase outlined above), these {\\it do not exist} in the standard Shu theory. As we will see, in practice both situations probably occur in nature. Observationally, one distinguishes various empirical stages in the evolution of young stellar objects (YSOs) from cloud core to (low-mass) main sequence star (e.g. Lada 1987, Andr\\'e 1994). The youngest observed YSOs are the Class~0 sources identified by Andr\\'e, Ward-Thompson, \\& Barsony (1993 -- hereafter AWB93), which are characterized by very strong emission in the submillimeter continuum, virtually no emission below $\\lambda \\sim 10\\ \\mu$m, and powerful jet-like outflows. Their very high ratio of submillimeter to bolometric luminosity suggests they have $M_{env} >> M_{\\star}$. Thus, Class~0 YSOs are excellent candidates for being very young protostars (estimated age $\\sim 10^4$~yr) in which the hydrostatic core has formed but not yet accreted the bulk of its final mass (AWB93). The next, still deeply embedded, YSO stage corresponds to the Class~I sources of Lada (1987), which are detected in the near-infrared ($\\lambda \\sim 2\\ \\mu$m) and have only moderate submillimeter continuum emission (Andr\\'e \\& Montmerle 1994; hereafter AM94). They are interpreted as more evolved protostars (typical age $\\sim$ 10$^5$~yr) surrounded by both a disk and a residual circumstellar envelope of substellar mass ($\\sim$ 0.1-0.3~$\\sm$ at most in $\\rho$ Ophiuchi; cf. AM94). Finally, the most evolved (Class~II and Class~III) YSO stages correspond to pre-main sequence stars (e.g., T~Tauri stars) surrounded by a circumstellar disk (optically thick and optically thin at $\\lambda \\la 10\\ \\mu$m, respectively), but lacking a dense circumstellar envelope.\\\\ Many examples of pre-collapse, pre-stellar cloud cores are known. In particular, Myers and co-workers have studied a large number of ammonia dense cores without $IRAS$ sources (e.g. Benson \\& Myers 1989) which are traditionally associated with (future) sites of {\\it isolated} low-mass star formation (see Myers 1994 for a recent review). These starless dense cores, which are gravitationally bound and close to virial equilibrium, are believed to be magnetically supported and to progressively evolve towards higher degrees of central concentration through ambipolar diffusion (e.g. Mouschovias 1991).\\\\ More compact starless condensations have been identified in regions of {\\it multiple} star formation such as the $\\rho$ Ophiuchi main cloud (e.g. Loren, Wootten, \\& Wilking 1990, Mezger et al. 1992b, AWB93). In these regions, the weakness of the static magnetic field (Troland et al. 1996) and the complexity of its geometry (Goodman \\& Heiles 1994 and references therein) suggest that gravitational forces overpower magnetic ones and that an external trigger rather than ambipolar diffusion is responsible for cloud fragmentation and core formation (e.g. Loren \\& Wootten 1986). In the present paper, we bring together two independent sets of observational results recently obtained on the density structure of pre-stellar dense cores (Sect. ~2.1) and on the evolution of protostellar outflows (Sect. ~2.2), and interpret them by means of a simple analytical theory (Sect. ~3.3) which sheds light on the accretion histories found in the numerical work of Foster \\& Chevalier (1993 -- hereafter FC93). We compare our theoretical predictions with observations in Sect.~4 and we conclude in Sect. ~5. ", + "conclusions": "The main points of our paper may be summarized as follows: \\begin{enumerate} \\item The radial density gradient of pre-stellar cores is typically flatter than $\\rho(r) \\propto r^{-1}$ near their centers and approach $\\rho(r) \\propto r^{-2}$ only beyond a few thousand AU (WSHA, AWM96, Fig.~1). In some cases, sharp outer boundaries, much steeper than $\\rho(r) \\propto r^{-2}$, are observed at a finite core radius (e.g. Abergel et al. 1996). This raises the possibility that, in some instances at least, the initial conditions for protostellar collapse depart significantly from a singular isothermal sphere. \\item The outflow momentum flux of embedded protostellar sources correlates very well with their circumstellar envelope mass (see Fig.~2 and Fig.~7). This suggests that the mass ejection and mass accretion rates of protostars both decline with time during protostellar evolution from Class~0 to Class~I sources (BATC). \\item Recent self-consistent hydrodynamical calculations of protostellar collapse indicate that initial conditions characterized by centrally flattened density profiles lead to a transitory phase of enhanced accretion immediately following the formation of the central hydrostatic protostar. This behavior is found regardless of whether the influence of magnetic fields is ignored (FC93) or fully accounted for (Tomisaka 1996). The nature of the transitory accretion peak is, however, left somewhat unclear by these numerical simulations. \\item On the basis of (3), we propose that (2) is a direct consequence of (1). \\item In order to elucidate the physical origin of (3), we analytically follow the history of protostellar accretion, using Lagrangian calculations based on a new type of self-similar `gravity-dominated' collapse solutions (Sect.~3.2, Sect.~3.3, and Appendix~A). We claim that these new solutions provide a better description of the collapse than the usual `pressure-dominated' self-similar solutions (e.g. Shu77, WS85), when pressure becomes negligible in the self-gravitating flow sometime after the onset of collapse, prior to stellar core formation. Comparison with the numerical simulations shows that the gravity-dominated solutions are adequate, at least qualitatively, in the supersonic region. \\item In agreement with (1), we start our supersonic calculations at $t = t_o < 0$ from an idealized pre-stellar core consisting of a strictly flat inner plateau up to a radius $r_N$, an $r^{-2}$ `envelope' up to $r_b$, and a steeper power-law `environment' farther out (Fig.~3). In our gravity-dominated description, the central plateau region first collapses homologously to form a {\\it finite-mass} hydrostatic stellar core at $t = 0$. Observationally, this initial phase, which does not exist in the standard Shu picture, should correspond to `isothermal protostars', i.e., supersonically collapsing cloud fragments with no central YSOs. It is followed at $t > 0$ by the main accretion/ejection phase, during which the non-zero central point mass accretes the surrounding envelope. As long as the gravitational influence of the initial point mass is significant, the accretion rate remains higher than the Shu value, $a_{eff}^3/G$. It then quickly converges towards $a_{eff}^3/G$. At late times, accretion of the outer environment leads to a terminal phase of residual accretion/ejection, during which the accretion rate declines below the Shu value (see Fig.~5). \\item We use our analytical model to fit the diagram of outflow efficiency $\\fco \\, c /\\lbol$ versus normalized envelope mass $\\menv / \\lbol^{0.6} $ obtained by BATC (see Fig.~7). To the extent that there is a direct proportionality between accretion and ejection, this diagram should provide an empirical measure of the accretion history of the sampled protostellar objects. A good overall fit is found when the model boundary radius $r_b$, is not much larger than the radius $r_N$ of the flat inner plateau, i.e., $r_b \\sim 1.6\\, r_N$. This requires that the fraction of cloud mass in the central plateau region be relatively large, $M_N/M_{cloud} \\sim 30$~\\%. (However, since the collapse is in fact less violent than in our pressure-less calculations, these values of $r_b$ and $M_N$ should be taken as indicative, and are likely to represent only a lower limit and an upper limit, respectively.) \\item Based on (7), we tentatively associate the short period of energetic accretion/ejection predicted by our model at the beginning of the accretion phase with the observationally-defined Class~0 stage (Sect.~4.1 and Fig.~7). In this view, Class~I objects are more evolved and correspond to the longer period of moderate accretion/ejection when the accretion rate approaches the Shu value. (See, however, point 10 below.) \\item We also find that the Ophiuchus $\\fco \\, c /\\lbol$ versus $\\menv / \\lbol^{0.6} $ diagram differs markedly from the corresponding diagram in Taurus: a clear contrast between Class~0 and Class~I objects is observed in Ophiuchus (Fig.~9) which is not seen in Taurus (Fig.~10). This points to different accretion histories in these two nearby star-forming clouds, which we interpret as arising from differences in initial conditions (Sect.~4.3). Both outflow and dense core observations suggest that the relative importance of the central plateau region in the initial density profile is significantly larger in Ophiuchus ($r_N \\sim r_b $, $M_N/M_{cloud} \\simgt 30$~\\% according to our approximate model) than in Taurus ($r_N \\sim 0.2\\, r_b$, $M_N/M_{cloud} < 10$~\\%). \\item According to our model fit of the Ophiuchus $\\fco \\, c /\\lbol$ versus $\\menv / \\lbol^{0.6} $ diagram, most of the $\\rho$~Oph Class~I YSOs should be in their terminal accretion phase (see Fig.~9). This is not surprising since the relatively small fragmentation length scale observed in $\\rho$~Oph implies that only a finite reservoir of mass is effectively available for the formation of any given protostar. \\item In conclusion, the `standard' theory of Shu and co-workers appears to describe protostellar evolution quite satisfactorily in regions of isolated star formation like Taurus. In our view, this is because in these regions most stars probably form following the {\\it self-initiated} contraction/collapse of dense cores due to ambipolar diffusion.\\\\ However, in star-forming clusters such as $\\rho$~Ophiuchi, the standard Shu theory is less appopriate since it does not account for fragmentation and multiple star formation. In these regions, star formation may be {\\it induced} by the impact of (slow) shock waves (e.g. Boss 1995), and protostellar cores may form by supersonic (or superalfv\\'enic) implosion of dense clumps, rather than slow (subsonic) ambipolar evolution. In this case, the collapse/accretion history advocated in the present paper (Sect.~3.3) and summarized in (6) is likely to provide a better description (see Fig.~9 and Sect.~4.3). \\end{enumerate}" + }, + "9701/astro-ph9701162_arXiv.txt": { + "abstract": "s{ Although they were discovered more than 25 years ago, gamma-ray bursts are still a mystery. Even their characteristic distance is highly uncertain. All that we can be confident about is that they involve compact objects and relativistic plasma. Current ideas and prospects are briefly reviewed. There are, fortunately, several feasible types of observation that could soon clarify the issues.} ", + "introduction": " ", + "conclusions": "" + }, + "9701/astro-ph9701212_arXiv.txt": { + "abstract": "We have simulated disk galaxies undergoing continual bombardment by other galaxies in a rich cluster. ``Galaxy harassment\" leads to dramatic evolution of smaller disk galaxies and provides an extremely effective mechanism to fuel a central quasar. Within a few billion years after a small disk galaxy enters the cluster environment, up to 90\\% of its gas can be driven into the inner 500 pc. Up to half of the mass can be transferred in a burst lasting just 100-200 Myr. This transport of gas to the center of galaxy is far more efficient than any mechanism proposed before. Galaxy harassment was first proposed to explain the disturbed blue galaxies in clusters seen in clusters at ($z \\gsim 0.3$), the ``Butcher--Oemler effect\". Quasars at the same reshifts lie in more clustered environments than those at lower redshift. Recent HST observations find that roughly half of all observed quasar host galaxiess are fainter than \\l*, with many of these less luminous hosts occuring at redshifts $z \\gsim 0.3$. We examine 5 quasars that are claimed to have low luminosity hosts and find that 3 are in rich clusters of galaxies, the fourth may be in a cluster but the evidence for this is marginal. The environment of the fifth has not been studied. ", + "introduction": "To power the central black hole by accretion, bright quasars must consume $10^8-10^{9}$ solar masses during their lifetimes---roughly 1\\% of the stellar mass in a bright elliptical galaxy or 10\\% of the gaseous mass of a bright spiral. Even with the assumption of a large host galaxy, an efficient mechanism is needed to channel gas to the central source. This problem of fuelling quasars is normally divided into three parts: the movement of gas from galactic scales to the inner few hundred parsecs, the instabilities of a self gravitating disk that transports the gas to a compact accretion disk and the detailed dynamics and radiation mechanisms of the final accretion. Here, we consider the first problem of moving gas from galactic scales into the inner few hundred parsecs. At redshifts of 2 where the number densities of quasars peak, this has been associated with the dynamical chaos of galaxy formation (c.f. Haehelt and Rees 1993). At lower redshifts, interactions of galaxies leading to coalescence or merging have been proposed (Hernquist 1989). These mergers can drive more than 10\\% of the gas to the center. The final host galaxy is the product of a merger, therefore it would be brighter than average. There are two observations that suggest a third mechanism at intermediate redshifts $0.2 \\lsim z \\lsim 0.8$. The environment of quasars has been observed to change with redshift (Yates, Miller and Peacock 1989; hereafter YMP89, Yee and Ellingson 1993; hereafter YE93). Quasars at higher redshifts are in Abell richness class 0-1 clusters of galaxies---an environment that is considerably richer than that of lower redshift quasars (the break point is at $z \\sim 0.3$ in YMP89 and $z \\sim 0.6$ in YE93). Bahcall, Kirhakos and Schneider (1995; hereafter BKS) used HST to image eight luminous quasars at redshifts between 0.15 and 0.3. Only three quasars in their sample have candidate hosts that are as luminous as {\\l*}, the characteristic luminosity in the Schechter (1976) luminosity function (LF). The other five hosts must be fainter than {\\l*} to have escaped detection. After observations of additional quasars, Bahcall {et al.} (1997; hereafter BKSS97) conclude that ``the luminous quasars studied in this paper occur preferentially in luminous galaxies\". They are rejecting the ``null hypothesis\" that all galaxies are equally likely to have quasars (e.g. a hypothesis that states that Draco and the Milky Way are equally likely to host quasars). Their conclusion results because at least half of all galaxies are $\\gsim 2$ magnitudes fainter than \\l* whereas the dividing line for their sample of quasar hosts is $\\sim${\\l*} within their errors. This depends slightly on the logarithmic slope of the LF. Since most LFs are weakly diverging at the faint end $N \\propto L^{-x}$, $1.5>x>1$, a faint end cutoff is required to to define an ``average luminosity\" and that average is normally 2 to 3 magnitudes brighter than the cutoff or $\\gsim 2$ magnitudes fainter than {\\l*}. However, galaxies brighter than $0.7-0.8${\\l*} contain half of all the luminosity, where the range takes into accounts the uncertainty in the faint end slope and cutoff of the luminosity function. This dividing line of luminosity is consistent with the BKSS97 midpoint of quasar hosts within their errors. So, quasars don't obviously prefer brighter galaxies any more than stars do. Previously, McLeaod and Rieke (1995) suggested a linear relation between the absolute magnitude of the quasar and its host. BKSS97 find no such relation other than what can be attributed to the obvious bias from detection limits. The simplest summary of the observations to date is that quasars and galaxies may be related much as stars and galaxies; the probability of finding either in a galaxy is proportional to the galaxies luminosity but their individual luminosities are not determined by the luminosity of their host. Hence, {\\it we need a mechanism that will produce quasars with a frequency per unit luminosity rather than a mechanism that only operates in bright galaxies as would be the case if mergers were the dominant trigger}. We have found a mechanism that is extremely efficient at channeling gas into the center of sub-\\l* galaxies that live in clusters. ``Galaxy harassment\" drives dynamical instabilities that send most of the gas into the central few hundred parsecs of the harassed galaxies. In \\S 2, we present detailed hydrodynamical simulations that show these effects, while we look at the model's predictions in \\S 3. ", + "conclusions": "Recent HST results have challenged the conventional picture that quasars are associated with high luminosity hosts. We have presented a model--- galaxy harassment---that rapidly channels gas into the centers of low luminosity hosts. This model has several predictions. We find evidence that most if not all of the quasars with sub-{\\l*} hosts are in the high density environments that lead to harassment. The HST images of host candidates by BKS95 shows tantalizing evidence of the distortions associated with harassment. In attempting to test the prediction of a higher frequency of AGNs in Butcher-Oemler clusters, we find some support amidst active controversy." + }, + "9701/astro-ph9701024_arXiv.txt": { + "abstract": "I consider the recent discovery of a soft X--ray source inside the error box of the gamma ray burst GB 960720 by the SAX, ASCA and ROSAT satellites, in terms of the fireball model. I show that the ejecta shell, which, after causing the burst is cold and dense, but still relativistic, keeps plowing through the interstellar medium, heating up the just--shocked matter which then emits X--rays. I compute the radiation emitted by this matter. I show that, up to about two months after the burst, in the cosmological scenario a soft X--ray ($0.1-10 \\; keV$) flux of at least $\\approx 10^{-13} \\; erg \\; s^{-1} \\; cm^{-2}$, well within current observational capabilities, is generated, explaining the observations of the three satellites. Instead, in the Galactic Halo scenario a flux $3$ orders of magnitude lower is expected. Detection of this non--thermal, declining flux in a statistically significant number of objects would simultaneously establish the fireball model and the cosmological nature of gamma ray bursts. ", + "introduction": "Our current theoretical understanding of gamma ray bursts (GRBs) is mostly based upon the fireball model (\\Mesz, Laguna, Rees 1993). However, though this model has won great critical acclaim because of its ability to explain two otherwise mysterious features (the bursts' time duration and the nonthermal spectra), it still has made no testable predictions that would allow gauging it against observations. This is mostly due to our ignorance of the physics of electron acceleration at relativistic shocks. To get a feeling of how serious this problem is, notice that the very same, simplified analysis of the radiation emitted behind the shock was carried out in two papers treating exactly the same physical problem, but separated by twenty years, Blandford and McKee (1977, dealing with AGNs) and Sari, Narayan, Piran (1996, dealing with GRBs), and that furthermore the first one was published before particle acceleration at shocks was even discovered (Bell 1978). It is the aim of this paper to derive a prediction from the fireball model, by (nearly completely) circumventing the problem of electron acceleration at relativistic shocks. Surprisingly, the prediction is different for cosmological and Galactic Halo scenarios. The opportunity to do this is offered by the detection, by the X--ray satellite SAX (Piro \\etal, 1995), of the gamma ray burst GB960720, both in the hard X--ray/soft $\\gamma$--ray band, where its results are confirmed by simultaneous observations by BATSE in a nearly identical band, and in the soft X--ray band (Piro \\etal, 1996a). Subsequent observations of GB 960720 (Piro \\etal, 1996b, Murakami \\etal, 1996, Greiner \\etal, 1996) made $\\approx 45$ days after the burst, have shown that a weak source, which is not an AGN, is present inside the WFCs' error box, $\\approx 5\\; arcmin$, with a flux $\\approx 2\\times 10^{-13}\\; erg\\; s^{-1}\\; cm^{-2}$ in the band $0.1-10 \\; keV$. It is thus interesting to speculate about what a systematic search for GBRs' afterglow ought to yield, in the fireball model. This afterglow has a simple interpretation: it is the cooling of matter swept up by the (still relativistic) shell of ejecta, plowing through the interstellar medium. In the next Section, I compute the expected soft X--ray fluxes from GRBs some time after the burst, for the fireball model, both in the cosmological and in the Galactic Halo scenarios. In the last Section, I discuss how the theoretical computations relate to the SAX, ASCA and ROSAT observations, and why the soft X--ray band is ideal for carrying out a statistical search for a soft X--ray afterglow. ", + "conclusions": "The most promising way to attack this problem is to try to identify afterglow emission from GRBs' error box in the soft X--ray on a statistical basis. The advantage of doing this in the soft X--ray rather than in lower energy bands, where source contamination is also low, is that the expected fluxes from Eq. \\ref{fx} can be seen to be at least an order of magnitude higher. Higher--energy bands, instead, do not have the angular resolution necessary to avoid source confusion. In the radio band, where lower fluxes are offset by much larger collecting areas, a behaviour similar to that of the soft X--ray band is of course expected, but theoretical computation is made difficult by several subtleties. Consideration of this effect is thus postponed to a forthcoming paper. The small angular resolution of the WFCs onboard SAX ($5\\; arcmin$) allows follow up observations by narrow field instruments with the hope of little source confusion. This is exactly what has been done by Greiner \\etal (1996), who identified three sources with the ROSAT HR Imager inside SAX error box. Of these, two sources are AGNs, while the third one, accounting for about half of the total flux (\\ie, $\\approx 2\\times 10^{-13} \\; erg\\; s^{-1} \\; cm^{-2}$) detected by SAX and ASCA, has no optical counterpart, indicating a very unusual object. From Fig. 1, where the tickmark indicates the position of $43$ post--burst days, and Eq. \\ref{flux2}, we see that the expected flux, $5\\times10^{-14} \\; erg\\; s^{-1}\\; cm^{-2}$, compares remarkably well (and perhaps fortuitously, since we do not know the source distance) with the observation of $10^{-13} \\;erg\\; s^{-1}\\; cm^{-2}$ (Piro \\etal, 1996b, Murakami \\etal, 1996, Greiner \\etal, 1996). Also, it should be noticed that Piro \\etal\\/(1996a) have set an upper limit to the soft X--ray flux from the burst region, immediately after the burst, of $10^{-10} \\; erg \\; s^{-1} \\; cm^{-2}$. Using Eq. \\ref{flux2} and Fig. 1 it can be seen that the highest flux expected in this model is $\\la 10^{-11} \\; erg \\; s^{-1} \\; cm^{-2}$. Thus, the current model reproduces correctly, for the most trivial choice of parameters, the observed features of the afterglow of GB960720. Hopefully, future observations ought to show that the afterglow has disappeared on a timescale of a few months, even though for this source it will be impossible to determine whether the light curve follows Fig. 1. Also, the apparent lack of interstellar absorption is consistent with the fireball model. The total baryon contamination $\\approx 10^{27} g$ (times at most a factor of $2$ to include the mass swept up when the shell is still relativistic) when spread out over a spherical surface of radius $R_{sh} < R < R_{rel}$ (Eqs. \\ref{rshock} and \\ref{rrel}) provides a total column depth $N_H \\approx 10^{16} - 10^{17} \\; cm^{-2}$, well within observational constraints. If furthermore, as has been suggested, GRBs are related to mergers of neutron star binaries, they are expected to be distributed somewhat like pulsars, \\ie\\/ outside the galactic disk, where column depths do not approach the observational bounds. Lastly, I would like to point out the reason why detection of soft X--rays about a month after the burst can discriminate between cosmological and Galactic Halo models, since it is rather amusing. In the fireball model, there are $5$ dimensional parameters: $E, M_{ej}, \\rho_{ISM}, c, D$, where $\\rho_{ISM}$ is the density of the circumstellar matter, and $D$ is source distance. Of these, $c$ is a universal constant, and $\\rho_{ISM}$ is an external parameter, which ought to be considered as given, and which is, furthermore, relatively well--known, compared with the uncertainty in parameters such as $E, M_{ej}, D$, each spanning several orders of magnitude. We can thus regard it as fixed. Thus, specifying that the two observational constraints, flux and time duration at Earth, be reproduced means fixing $2$ of the $3$ free parameters, leaving only one (say, $D$) undetermined. This corresponds to having a one--parameter ($D$) family of homologous solutions, each fitting observational data, each located at different distances from the observer. By making observations at a fixed time after the burst, about a month, we are observing cosmological and Galactic Halo models at non--homologous moments, thus breaking the similarity law that links them: in fact, the shock is still relativistic in the cosmological scenario, and well subrelativistic in the Galactic Halo scenario. In short, I have argued that, in the cosmological scenario of the fireball model, detectable fluxes (Eq. \\ref{flux2} and Fig. 1) of soft X--rays should be emitted in the two months following a gamma ray burst, with a non--thermal spectrum and a characteristic decrease (Fig. 1), while no detectable flux can arise in the Galactic Halo scenario. In particular, the expected flux of $10^{-13} \\; erg\\; s^{-1} \\; cm^{-2}$ compares remarkably well with the observations of GB960720 made about $40$ days after the burst, by Piro \\etal (1996b), Murakami \\etal, (1996) and Greiner \\etal (1996). But the model also predicts the time--dependence of the afterglow (Fig. 1) and its disappearance after a few months, and is thus subject to more elaborate testing. Detection of the afterglow in a statistically meaningful sample of gamma ray bursts would simultaneously establish both the fireball model {\\bf and} the cosmological nature of GRBs. \\vskip 1truecm I am indebted to L. Stella and especially to Luigi Piro, for fruitful scientific conversations, and to an anonymous referee for constructive criticisms." + }, + "9701/astro-ph9701030_arXiv.txt": { + "abstract": "Using the Faint Object Camera on the repaired {\\it Hubble Space Telescope,} we have observed two fields in the globular cluster M15:\\ the central density cusp, and a field at $r = 20''$. These are the highest-resolution images ever taken of this cluster's dense core, and are the first to probe the distribution of stars well below the main-sequence turnoff. After correction for incompleteness, we measure a logarithmic cusp slope ($d \\log \\sigma / d \\log r$) of $-0.70 \\pm 0.05$ (1-sigma) for turnoff ($\\sim0.8 \\Msun$) stars over the radial range from $0\\secspt3$ to $10''$; this slope is consistent with previous measurements. We also set an approximate upper limit of $\\sim1$\\secspt5 (90\\% confidence limit) on the size of any possible constant-surface-density core, but discuss uncertainties in this limit that arise from crowding corrections. We find that fainter stars in the cusp also have power-law density profiles:\\ a mass group near $0.7 \\Msun$ has a logarithmic slope of $-0.56 \\pm 0.05$ (1-sigma) over the radial range from $2''$ to $10''$. Taken together, the two slopes are not well matched by the simplest core-collapse or black-hole models. We also measure a mass function at $r = 20''$, outside of the central cusp. Both of the FOC fields show substantial mass segregation, when compared with a mass function measured with the WFPC2 at $r = 5'$. In comparing the overall mass functions of the two FOC fields and the $r=5'$ field, we find that the radial variation of the mass function is somewhat less than that predicted by a King--Michie model of the cluster, but greater than that predicted by a Fokker--Planck model taken from the literature. ", + "introduction": "For years, studies of M15 have guided our understanding of the dynamical evolution of globular clusters. Observations have shown that the cluster's stellar density rises in its center in a power-law cusp, with no sign of a ``flat,'' constant-surface-density core at the resolution attainable from the ground (Djorgovski \\& King\\markcite{djk84} 1984, Lugger\\markcite{lug87} \\etal\\ 1987). Most researchers have considered M15 to be in a state of core collapse, in which the cluster core's negative heat capacity has resulted in its runaway contraction. The collapse would eventually be halted by energy input from binary stars (see the references in Hut \\& Makino\\markcite{iau} 1996). Alternatively, M15 could harbor a massive ($\\sim10^3 \\Msun$) black hole in its center; the cusp would then be a result of the black hole's deep potential well (Bahcall \\& Wolf\\markcite{bw76} 1976). The increase in resolving power provided by the {\\it Hubble Space Telescope} ({\\it HST}) can help to determine the nature of the density cusp. The principal goal of \\hst\\ observations of M15 has been to resolve stars near the main-sequence turnoff, since the large number of such stars should trace the stellar density profile much more clearly than the handful of bright giants seen from the ground. Early \\hst\\ results, though, provided more confusion than clarity:\\ until the 1993 repair, most of the light of the bright stars was spread into an extended ``halo'' that obscured the view of the faint stars. Lauer \\etal\\ (1991)\\markcite{lau91} attempted to subtract these bright stars from a $U$-band Planetary Camera image of the M15 cusp, and claimed the detection of a ``flat'' (constant surface-density) core of radius 2\\secspt2 in the residual light. Yanny\\markcite{yan94} \\etal\\ (1994), however, found a serious flaw in this procedure:\\ stellar-photometry software attributes too much of the light to the bright stars in a crowded field, at the expense of their faint neighbors. The bright stars are thus oversubtracted in the densest region near the cluster center, creating the illusion of a flat core. Guhathakurta\\markcite{GYSB} \\etal\\ (1996, hereafter GYSB) followed up the Yanny \\etal\\ study by taking post-repair WFPC2 images of the M15 cusp, and found that the stellar density rises all the way into $r \\simeq$ 0\\secspt3. Moreover, GYSB found that the cusp surface density is well fit by a power law of slope $-0.82 \\pm 0.12$---quite close to the slope of $-0.75$ expected for stars near a massive black hole, but also consistent with the outcome of core collapse in the presence of dark objects of mass $\\geq 1 \\Msun$ (Cohn\\markcite{hc85} 1985, Grabhorn \\etal\\ \\markcite{gra92}1992). Our main goal in this paper is to present an analysis of a set of images of the M15 cusp, taken with the Faint Object Camera (FOC) aboard the repaired \\hst. The FOC samples the \\hst\\ point-spread function (PSF) much better than the WFPC2:\\ its pixel scale is $0\\secspt014$/pixel, compared to the Planetary Camera chip's $0\\secspt044$/pixel (the FWHM of the $V$-band PSF is $0\\secspt04$). In the central $2''$ of M15, the typical separation of stars with $19 < V < 21$ is $\\sim 0\\secspt15$, so the FOC can separate many more close stellar pairs. We begin by describing the images themselves. We then describe our procedure for extracting as much photometric information as possible from the images, and the extensive artificial-star experiments that accompany this procedure. We then use these data to construct the surface-density profiles and mass functions of M15. Finally, we construct a simple dynamical model of M15, and compare our results with a Fokker--Planck model from the literature. ", + "conclusions": "We have presented radial distributions and mass functions of main-sequence stars in two fields near the center of M15. The distribution of turnoff stars in the innermost $10''$ is similar to what has been found previously: we find a power-law density cusp, with logarithmic slope $-0.70 \\pm 0.05$ (for stars with $r>0\\secspt3$). The density continues to rise in to small radii; any constant-surface-density core must be smaller than $1\\secspt5.$ We have shown, for the first time, that fainter, lower-mass stars also have a power-law distribution near the cluster center; we find a logarithmic slope of $-0.56 \\pm 0.03$ for a group of stars with masses near $0.7\\Msun$, over the radial range from $2''$ to $10''$. (Over the same range, turnoff-mass stars have a slope of $-0.64 \\pm 0.08.$) These slopes are not well matched by the predictions of the simplest core-collapse and black-hole models. We have also compared the mass functions in our fields with the MF in a field $5'$ from the cluster center, and found a strong degree of mass segregation. Finally, we fit a King--Michie model to the cluster's surface-density profile and outer mass function, and showed that it predicts somewhat less segregation than we observe. The model also does not match the cluster's velocity-dispersion profile. While our observations alone cannot settle the debate over the dynamical state of M15, we hope our results will provide a valuable set of inputs to a new generation of sophisticated and realistic models of globular-cluster evolution." + }, + "9701/astro-ph9701206_arXiv.txt": { + "abstract": "We give a global analysis of mass transfer variations in low-mass X-ray binaries and cataclysmic variables whose evolution is driven by the nuclear expansion of the secondary star. We show that limit cycles caused by irradiation of the secondary by the accreting primary are possible in a large class of these binaries. In the high state the companion transfers a large fraction of its envelope mass on a thermal timescale. In most cases this implies super--Eddington transfer rates, and would thus probably lead to common--envelope evolution and the formation of an ultrashort--period binary. Observed systems with (sub)giant secondaries stabilize themselves against this possibility either by being transient, or by shielding the secondary from irradiation in some way. ", + "introduction": "Semidetached binaries in which a compact object (white dwarf, neutron star or black hole) accretes material via Roche lobe overflow from a companion on or near the main sequence are of great interest in current astrophysics. The evolution of such systems is driven by orbital angular momentum losses via gravitational radiation and magnetic braking (see e.g. King 1988 for a review). Thus most properties of the binary, particularly the mean mass transfer rate, depend essentially only on the secondary mass $M_2$ and change on the timescale $t_J \\sim 10^8 - 10^9$ yr for angular momentum loss. In a recent paper (King, Frank, Kolb \\& Ritter 1996a, henceforth Paper I) we discussed the conditions under which mass transfer can vary cyclically about the evolutionary mean in these systems by developing a general formalism allowing one to study the stability of mass transfer in systems driven by angular momentum losses. The existence of cycles is required to account for the wide dispersion of e.g. mass transfer rates at a given orbital period. We concluded that the most likely cause of such cycles is weak irradiation of the companion star by the accreting component. The conditions for this appear to be fulfilled in a large class of cataclysmic variables (CVs), in which the accreting star is a white dwarf. Of course semidetached evolution is not restricted to systems with main--sequence companions. In systems with orbital periods $P\\gta 1$ day the orbital evolution and hence the rate of mass transfer is either determined or strongly influenced by nuclear evolution of the companion. In this paper we discuss the stability of mass transfer in systems containing a giant or a subgiant companion and consider the possibility of irradiation--driven mass transfer cycles similar to those thought to exist in CVs. For this purpose we generalize the analysis of Paper I to include the effects of nuclear evolution on the radius variations by using a simple core--envelope model for the companion. This is a good representation of systems with low--mass giant secondaries, which constitute the great majority of long--period compact binaries. However this description does not apply to the recently--discovered black--hole transient GRO J1655-40, where the companion star appears to be crossing the Hertzsprung gap (Orosz \\& Bailyn, 1996). ", + "conclusions": "We have shown that irradiation of an evolved low--mass companion in LMXBs and CVs can drive mass transfer cycles. These cycles do not need the intervention of any further effect, unlike the case of main--sequence companions (Paper I), where a modest increase in the driving angular momentum loss rate is required in the high state if cycles are to occur in many cases. The cycles with evolved companions are also considerably more violent than in the main--sequence case. This is a direct consequence of two facts. First, the nuclear luminosity of an evolved star is insensitive to the stellar radius, so that blocking of the intrinsic stellar flux by irradiation requires the star to expand so as to maintain the same unblocked area. This leads to much larger expansions than in the main--sequence case, where the nuclear luminosity drops sharply as the star expands. Second, the ratio $t_{\\rm dr}/t_{\\rm ce} = f\\rho$ of driving to thermal timescales is much larger in evolved stars than on the main sequence, making the expansion very rapid. In the high state an evolved star would lose a significant fraction of its total envelope mass on a thermal timescale. In a CV with a low core mass, the implied accretion rate would probably turn the system into a supersoft X--ray source. In CVs with higher core masses, the nuclear burning causes the white dwarf to develop an extensive envelope, while in long--period LMXBs, the high state accretion rates greatly exceed the Eddington limit. If the instability is not quenched all these systems would undergo a common--envelope phase. They may merge entirely, or reappear as ultrashort--period systems like AM CVn (for white dwarf primaries), or helium--star LMXBs, or detached systems with low--mass white dwarf companions. However, there are at least two ways in which the instability can be quenched and the systems (particularly LMXBs) stabilized: shielding of the companion by e.g. an extensive accretion disc corona, and intermittent accretion. The typical duty cycles $d \\lta 10^{-2}$ observed in soft X--ray transients are short enough to stabilize them, while observed dwarf nova duty cycles are unable to stabilize CVs with evolved companions. Although both means of stabilization seem to occur in nature, it is clear that the irradiation instability is so violent that it must play a major role in any discussion of the evolution of CVs and LMXBs with evolved companions: the systems must somehow stave it off or evolve catastrophically. We shall consider some of the observational consequences in a future paper. JF thanks the Max--Planck--Institut f\\\"ur Astrophysik for warm hospitality during a productive stay in June--August 1996. This work was partially supported by the U.K. Science and Engineering Research Council (now PPARC) and by NASA grant NAG5-2777 to LSU. ARK acknowledges support as a PPARC Senior Fellow, and the warm hospitality of the MPA. \\clearpage" + }, + "9701/astro-ph9701083_arXiv.txt": { + "abstract": "Rayleigh-Taylor (R-T) instabilities in the explosion of SN 1993J are investigated by means of two-dimensional hydrodynamical simulations. It is found that the extent of mixing is sensitive to the progenitor's core mass and the envelope mass. Because the helium core mass (3 - 4 \\ms) is smaller than that of SN 1987A, R-T instabilities at the He/C+O interfaces develop to induce a large scale mixing in the helium core, while the instability is relatively weak at the H/He interface due to the small envelope mass. The predicted abundance distribution, in particular the amount of the \\ni~ mixing, is compared with those required in the theoretical light curves and the late time optical spectra. This enables us to specify the progenitor of SN 1993J in some detail. ", + "introduction": "SN 1993J has been identified as a Type IIb supernova (SN IIb) from the spectral changes which show growing features of helium and oxygen, and from the optical light curve which shows double peaks (Wheeler \\& Filippenko 1996; Baron \\etal 1995 for reviews and references therein). These features are distinct from those of previously known Type II supernovae (SNe II). It was obvious that the peculiar light curve of SN 1993J cannot be accounted for by an explosion of an ordinary red supergiant with a massive hydrogen-rich envelope, which produces a light curve of a SN II-P; instead, it can be well reproduced as the explosion of a red-supergiant whose hydrogen-rich envelope is as small as \\(< \\) \\ 1 \\ms~ (Nomoto \\etal 1993; Podsiadlowski \\etal 1993; Shigeyama \\etal 1994; Bartunov \\etal 1994; Utrobin 1994; Woosley \\etal 1994; Young \\etal 1995). The progenitor of SN 1993J is likely to have lost most of its H-rich envelope due to the interaction with its companion star in a binary system. Constraints on mixing in the ejecta of SN 1993J come from both observed spectra and photometry. Spectroscopically, the asymmetric profiles of the [O I] and [Mg I] emission lines at late times indicate mixing in the ejecta (Spyromilio 1994; Wang \\& Hu 1994). Shigeyama \\etal (1994) first examined two extreme cases of mixing in their light curve modeling (i.e., with complete homogeneous mixing inside the helium layers and without mixing) and favored mixing of \\ni~ in SN 1993J. Woosley \\etal (1994) also favored mixing of \\ni~ in their light curve modeling. The occurrence of mixing and clumpiness in the supernova ejecta was first confirmed in SN 1987A. One clear confirmation is the early detection of hard X-rays and \\(\\gamma\\)-rays originated from the decays of \\ni~ and \\co~ (e.g., Kumagai \\etal 1989). Stimulated by these observations and theory, a number of multidimensional simulations have been performed to show that the R-T instability develops indeed in the ejecta of SN 1987A (Arnett \\etal 1989; Hachisu \\etal 1990, 1992; M\\\"uller \\etal 1991; Den \\etal 1990). It was found that the instability is weak at the He/C+O and Ni/O interfaces but strong at the H/He interface, which can be expected from the large helium core mass \\( \\sim \\) 6 \\ms~ and the massive hydrogen-rich envelope \\( \\sim 10 \\) \\ms. In SN 1987A, the large scale mixing of nickel was induced by falling spikes of hydrogen which penetrated into deep inner layers. Later, it has been recognized that the occurrence of mixing and clumping in the supernova ejecta can occur also in other types of supernovae and that the effect on their light curves can provide interesting diagnostics of the internal structure of the progenitors. The extent of mixing and clumpiness affect the light curve and spectra. A number of hydrodynamical simulations have shown that the R-T instabilities arise in the Type II-P explosions of red-supergiants (Herant \\& Woosley 1994; Shigeyama \\etal 1996) and helium star models of Type Ib supernovae (Hachisu \\etal 1991, 1994) as in SN 1987A. These studies indicate the importance of multidimensional simulations of instabilities in supernova ejecta. This is particularly interesting for SN 1993J because its presupernova structure is suggested to be quite different from other supernovae. Conversely, we may strongly constrain the progenitor's mass and structure because the R-T instabilities are sensitive to the density structure of the progenitor. Despite such importance, no such simulation has been conducted before. Therefore, we have carried out two dimensional hydrodynamical simulations for the R-T instability in the ejecta of SN 1993J to examine the extent of mixing and the resultant abundance distributions quantitatively. In parallel, we have calculated bolometric light curves for several explosion models with various extent of mixing and compared them with the observations. These results would enable us to constrain the progenitor model, which would be useful to understand the still debated presupernova evolution. In the next section, the presupernova models for SN 1993J are described. Results of the linear stability analysis and the 2D hydrodynamical simulations are shown in \\S3 and \\S4, respectively. The effects of \\ni~ mixing on the optical light curve are discussed in \\S5. Finally we summarize some constraints on the explosion models. ", + "conclusions": "We have investigated the Rayleigh-Taylor instabilities in the ejecta of SN 1993J with a linear analysis of spherically symmetric explosion models and with a two-dimensional hydrodynamical simulations. We find the following conclusions. 1. The instability at the He/C+O interface develops to induce a large scale mixing because of the relatively small He core mass. 2. The instability at the H/He interface is weak because of the small hydrogen-rich envelope mass. These features (1,2) are in contrast to SN 1987A which had the more massive He core and the envelope. The extent of mixing of heavy elements (Ni and C+O) is sensitive to the core mass. For the smaller core mass, the R-T instability is stronger and causes more extensive mixing due to the smaller mass ratio between the core and the He layer. 3. The optical light curves are calculated with a parameterized degree of mixing. The oberved light curve is well reproduced if substantial amount of \\ni~ mixing occurs. 4. The calculated abundance distributions of the ejecta against the expansion velocity are compared with the observed velocities of Ni, O, and H. The model with the 3.3 \\ms~ He core and the hydrogen-rich envelope of 0.3-0.4 \\ms~ can well reproduce the observational feature of SN 1993J, if the explosion energy is as low as $\\sim$ 0.8 \\e{51} ergs. The model with the 4 \\ms~ He core and the $\\sim$ 0.5 \\ms~ envelope is also a good alternative, if \\ni~ is more extensively mixed than our present calculations, possibly due to much larger initial perturbations. \\bigskip This work has been supported in part by the grant-in-Aid for Scientific Research (05242102, 06233101) and COE research (07CE2002) of the Ministry of Education, Science, and Culture in Japan, and the fellowship of the Japan Society for the Promotion of Science for Japanese Junior Scientists (4227). The computation was carried out on Fujitsu VPP-500 at the Institute of Physical and Chemical Research (RIKEN) and the Institute of Space an Astronautical Science (ISAS), and VPP-300 at the National Astronomical Observatory in Japan (NAO, Tokyo)." + }, + "9701/astro-ph9701227_arXiv.txt": { + "abstract": "The problem of automated separation of stars and galaxies on photographic plates is revisited with two goals in mind : First, to separate galaxies from everything else (as opposed to most previous work, in which galaxies were lumped together with all other non-stellar images). And second, to search optically for galaxies at low Galactic latitudes (an area that has been largely avoided in the past). This paper demonstrates how an artificial neural network can be trained to achieve both goals on Schmidt plates of the Digitised Sky Survey. Here I present the method while its application to large numbers of plates is deferred to a later paper. Analysis is also provided of the way in which the network operates and the results are used to counter claims that it is a complicated and incomprehensible tool. ", + "introduction": "The separation of stars from galaxies on photographic plates or CCDs is required in large surveys like the APM survey (Maddox {\\it et al.} 1990a) and the Sloan Digital Sky survey (Gunn 1995), for purposes such as the preparation of target lists of galaxies for observations (e.g., in order to obtain redshifts). The sheer number of detected objects on survey plates forces one to use automated procedures for this task and in every large survey some such method is chosen. It is generally relatively easy to separate single images of stars from everything else, and different methods have been successfully tried in order to achieve this goal. The DAOFIND package (e.g., as incorporated in IRAF) fits point spread functions to detect stars. Sebok (1979) uses surface brightness measurements from the image and matches them to templates of stars and various galaxy models. Jarvis and Tyson (1981) define shape parameters on the basis of image moments. Slezak {\\it et al.} (1988) combine light profile with shape parameters, while the APM survey team (Maddox {\\it et al.} 1990a) use automated parameters derived from the APM microdensitometer. Regardless of which parameters are measured, the problem is usually formulated in terms of finding an optimal decision surface in the space spanned by the chosen parameters. This approach implies supervised learning by the classifier, i.e., learning from examples that were pre-classified, e.g., by eye. An alternative approach was recently proposed by M\\\"{a}h\\\"{o}nen \\& Hakala (1995), who use an unsupervised learning method (self organising maps) to separate images on the basis of their appearance without classifying them by eye first. While promising, their approach was only demonstrated on synthetic images and is yet to be applied to real images. In this paper supervised learning was preferred. The typical output of a classifier is the drawing of linear boundaries which distinguish between stars and non-stellar images, and it is usually assumed that the vast majority of all non-stellar objects in the field are galaxies. There are at least three points that require further consideration when constructing yet another star/galaxy classifier : \\begin{enumerate} \\item{Which parameters are adequate for describing galaxies and stars ?} \\item{Is linear classification good enough ?} \\item{Is it true that non-stellar objects are mostly galaxies ?} \\end{enumerate} \\bigskip Although important, the first question is difficult to answer. Different parameters are measured by different researchers who do not all use the same plate material or even the same classification tools. Since the end result is a combination of all of these factors it is difficult to examine the effect of parameter choice separately. Because of this difficulty this paper makes no attempt to compare the parameters used here with those chosen by others. Rather, I augment parameters which were used successfully elsewhere (Odewahn {\\it et al.} 1992) by several special-purpose parameters. The second question requires careful comparison of linear and non-linear classifiers, using the same parameters and the same plate material. In many cases linear separation of stellar from non-stellar objects in the selected parameter space appears adequate (as judged by the distribution of galaxies and stars in plots of one parameter against another). However, in general a non-linear classifier is capable of at least as good a classification as a linear classifier, and thus the choice of a non-linear classifier seems natural. This paper follows Odewahn {\\it et al.} in selecting artificial neural networks as the (non-linear) classifier. As for the third question, it is shown below that galaxies do not dominate the population of non-stellar objects : the incidence of galaxies among non-stellar images is less than $50\\%$ even at high Galactic latitudes. This implies that treating ``non-stars'' as ``galaxies'' is in general wrong. Among the many non-stellar images on a typical plate one finds images of two or more overlapping stars which create elongated shapes; plate defects (e.g., scratches); meteorite trails; and combinations of the above with true galaxies. The fraction of galaxies drops significantly as one moves to low Galactic latitudes, making the problem worse still. The aim of this paper is to extend previous work in two respects : first, by constructing a classifier that is capable of telling galaxies from all other objects. Discriminating between merged and extended objects is not an easy task and is largely ignored in this field. A notable exception is the work done by the APM team (Maddox {\\it et al.} 1990a), who constructed two special parameters for telling merged objects apart from single objects. I attempt to carry this distinction further, by training the ANN to tell apart six different classes of object. The second respect in which previous work is extended is by making the classifier general enough to handle star/galaxy separation at low Galactic latitudes. The reason why one would wish to extend high latitude work into the {\\it zone of avoidance} is the difficulty of using standard techniques to detect galaxies optically behind the disk of the Galaxy. The potential for new discoveries in that part of the sky has been realised several times in the past few decades, most recently in the radio band by Kraan-Korteweg {\\it et al.} (1994). An optical identification of galaxy candidates could help save a lot of time in the search for partially obscured galaxies. This paper serves to introduce the problem and the chosen method. The ultimate goal is the production of an automated catalogue of galaxy candidates behind the Galactic disk. The catalogue will be described in a later paper, which will also address issues of data quality, plate to plate changes and the overall reliability of the classifications. The data used for this first paper are described in \\S~2. \\S~3 describes the eyeball classification of the images. The parameters chosen to represent each image are disussed in \\S~4 and put to use in training the ANN, in \\S~5. The discussion follows in \\S~6. \\bigskip ", + "conclusions": "Automated star/galaxy separation is not a new application in astronomy, but surprisingly work to date has left much to be desired. Separating stars from everything else is not enough for researchers whose main interest is galaxies, and we show that galaxies can be reliably separated from other non-stellar images at high Galactic latitudes. However, the potential for exciting discoveries is naturally larger in areas that were traditionally avoided, such as fields at low Galactic latitudes. We show that it is possible to perform the separation of galaxies from other objects in the Zone of Avoidance, although the success rates are lower and the contamination higher. The tool used here in the role of automated classifier is an ANN. Other classifiers exist, of course, and should give similar results with our data and its parametrisation. The choice of a classifier is largely a matter of convenience. Most of the hard work goes into the parametrisation of the problem. Nevertheless, ANNs have probabilistic capabilities that are useful, e.g., in trying to improve detection limits, and they are easy to use and very versatile. It has been claimed in the past that ANNs are slow, complicated and are difficult to understand or interpret. The ANN code used here, which was kindly supplied by B. Ripley, takes few CPU minutes to converge on a conventional workstation and is very easy to implement. The analysis of the ANN weights (\\S~5.4 above) shows that ANNs can be understood and provide insight as to which parameters are more important than others. The methods described here are readily applicable to any number of fields extracted from the DSS. It is quite possible to proceed and catalogue the entire sky in this manner. However, it might be of limited interest to proceed with the high Galactic latitude patches, as the available digitised plates do not go very deep and the sky has been mapped at these latitudes and down to the chosen limiting size. However, it is very interesting and apparently feasible to produce such a catalogue for the Zone of Avoidance, from which target lists for pointed observations of galaxies could be easily prepared. This project will be pursued further once the higher resolution DSS mk II becomes available. \\bigskip {\\bf ACKNOWLEDGEMENTS} \\bigskip The work described in this paper constituted one part of my Ph.D. thesis, under the inspirational supervision of Ofer Lahav. I am indebted to him for much insight into ANNs and their statistical nature. I am enormously grateful to Steve Maddox and Steve Odewahn for sharing their experience with me and bearing with me through long discussions. I would like to thank Brian Ripley for allowing me to use his ANN code, and to acknowledge an Isaac Newton studentship which supported me through part of this research. \\newpage \\bigskip Figure Captions : \\bigskip Figure 1 : A 15'x15' Section of one of the Patches taken from Field 646. Compare with figure 2. \\bigskip Figure 2 : A 15'x15' Section of one of the Low Galactic Latitude Patches. Compare with figure 1. \\bigskip Figure 3 : Distributions of Eyeball-Classes in the High Latitude and Low Latitude Datasets. Note the total dominance of stellar mergers at low Galactic latitudes. \\bigskip Figure 4 : Weights Connecting Inputs to Hidden Layer Nodes for a Single Run. \\bigskip Figure 5 : Weights Connecting Hidden Layer Nodes to Outputs for a Single Run. \\bigskip" + }, + "9701/astro-ph9701157_arXiv.txt": { + "abstract": "Since the discovery of the Ultraviolet Upturn Phenomenon (``UVX'') in early-type galaxies it has been clear that the stellar populations of such systems contain an unexpected hot component. Recent work has provided strong circumstantial evidence that the stars radiating at short wavelengths (\\wless{2000}) is in fact due to hot horizontal branch, post-HB stars and post-AGB stars. We summarize the arguments in favour of this hypothesis. We then derive an estimate for the fraction of all HB stars that must be contributing to the UV upturn phenomenon in the strongest UVX galaxy, NGC 1399, and derive a hot star fraction $f_{H} \\sim 0.16.$ The implication is that UVX arises from a minority fraction of the dominant stellar population. We conclude that the mechanism that produces the UVX is not one that can be explained naturally by the presence of an extremely metal-rich or metal-poor population. ", + "introduction": "The Ultraviolet Upturn Phenomenon was first found by the OAO-2 spacecraft in the late sixties (\\cite{co69}). It consists of a `UV rising branch' at wavelengths shorter than \\wless{2500}, which varies in amplitude amongst the galaxies observed. The amplitude of the UVX phenomenon varies by a factor of about 10 from NGC 1399 [$\\xvv = 2.05$] to M32 [$\\xvv = 4.50$] with very similar slopes in the IUE spectral range (\\wless{1200}).The question was investigated in further detail by Burstein \\etal\\ (1988, hereafter B3FL) who presented a sample of galaxies observed by the International Ultraviolet Explorer (IUE). They determined the UV/Optical color $\\mlam(1550) - V$ [hereafter \\xvv] from ground-based photometry. They plotted these colors against the absorption line index \\mgii, a reliable metallicity indicator for the Galactic globular cluster system (\\cite{bfgk}; \\cite{bh90}), and the central velocity dispersion $\\sigma_0.$ For the galaxies that are quiescent, i.e. which contain no evidence of recent star formation and nuclear activity there is a non-linear correlation between \\mgii\\ and the UVX amplitude (see also \\cite{f83}). The UVX was also found to correlate, albeit not as tightly, with $\\sigma_0,$ which is an indicator of luminosity through the Faber-Jackson relation. Thus either the amplitude increase is directly related to the metallicity (or at least magnesium abundance) of the galaxies, or there is a tendency for bright galaxies to have strong UV upturns. Enhancements in $\\alpha$-capture elements such as Mg have been found to be a feature of the spectra of bright ellipticals (\\cite{wfg92}) and imply something about the early enrichment history of massive systems. This contribution is organized as follows: in the next section we discuss the recent observational evidence pertaining to the UVX. Section 3 summarizes the problems associated with modelling the Horizontal Branch. In \\S 4 we give an overview of evolutionary population synthesis, and use it to derive estimates of the fraction of hot stars present as a function of the UV upturn strength. \\S 5 discusses these estimates in the light of population models for the UVX. ", + "conclusions": "We have focussed here on deriving limits on the population size responsible for the UV upturn. This is an important prediction of the theory, since is is the most easily tested for consistency against other observational consequences. In particular, for the case of NGC 1399 [$\\xvv = 2.05$] we obtain $\\fuv \\sim 0.16.$ This conclusion implies that the UVX sources do not arise from a trace population more prevalent in UVX-strong galaxies, and instead favours the notion that a sizeable minority of the dominant population is contributing to the UVX. The work of \\cite{pl97} and \\cite{bcf94} have assumed that a fraction of stars at the low and high end of the metallicity distribution contributes the UV radiation from the galaxies. In the case of the metal-poor hypothesis, a 16\\% fraction of the population in stars with $\\rm [Fe/H] < -1$ would give a large enough contribution to the mid-ultraviolet radiation to produce a discrepancy with observation (\\cite{wdj96}; Bressan \\etal\\ 1994). Quantitative estimates employing the 2500 \\AA\\ fluxes in DOR95 in equation~(\\ref{final}) show that bimodal metallicity distributions, in which the majority population contributes relatively little in the mid-ultraviolet, are necessary to reconcile that longer wavelength flux with observations. The hypothesis that the UVX stars are entirely metal-rich clearly does not suffer from this problem. However, the large fraction of stars necessary to account for the observed flux seems to imply that a spread in metallicity must give rise to the UVX rather than the most extreme composition. This is contrary to the work of Bressan \\etal\\ who suggest that the UVX arises from the effects of stellar evolution at large ages. They invoke high metallicity, strongly helium enhanced models with RGB mass loss similar to what is inferred from the globular cluster system. Assuming that it is a fraction of the dominant population that gives rise to the UVX accords better with the observational record we have for Galactic populations. We do not need to hypothesise that {\\it all} RGB stars of a given metallicity have the same degree of mass loss, which we do not observe in any other context (little information about mass distributions can be deduced from red HB clumps). Unfortunately one cannot test the initial metallicity of the field sdB stars easily since their spectra are affected by diffusion in the surface layers. However, the question of how the UVX arises remains open, at least for the time being (but see \\cite{d97a} for empirical uses of the far-UV radiation). It should be stressed that both `metal-poor' and `metal-rich' models as presented by these authors rely on the hypothesis that age is established as a `global' second parameter for HB morphology in old stellar systems. They use the results of simple, unimodal synthetic HB models to derive ages of galaxies. We argue that the Galactic record in resolved populations does not support the use of these models and cannot therefore constrain unresolved stellar populations. We close with a caveat that applies to the interpretation of the UVX as a metallicity-driven phenomenon, on which the `metal-rich' hypothesis relies heavily. We caution that this assumption should not be regarded as incontrovertible. The argument for a metallicity driven mechanism is weakened by the discovery that the \\mgii\\ indicator does not trace the heavy-element abundance in galaxies and by the observation that \\mgii\\ line strengths are correlated with $\\sigma_0$ but not with iron indicators (\\cite{wfg92}). The UVX may thus be uncorrelated with iron abundance (\\cite{d97a} and references therein). Another definite prediction of the metallicity relation is that the UVX gradient with galactocentric radius should be correlated with optical line-index gradients. This question is currently under study using Ultraviolet Imaging Telescope data (Ohl \\etal, in preparation). Thus, the Ultraviolet Upturn Phenomenon remains, like many others in astronomy, an unsolved problem which requires spatially resolved spaceborne observations for further study. Acknowledgements: B.D. acknowledges support from NASA grants NAG5-700 and NAGW-4106, and healthy discussions with Robert T. Rood and Robert W. O'Connell. He is also glad to report that Icko Iben Jr., who we honour with this volume, did not find anything wrong with his presentation." + }, + "9701/astro-ph9701139_arXiv.txt": { + "abstract": "The defining characteristic of a black hole is that it possesses an event horizon through which matter and energy can fall in but from which nothing escapes. Soft X-ray transients (SXTs), a class of X-ray binaries, appear to confirm this fundamental property of black holes. SXTs that are thought to contain accreting black holes display a large variation of luminosity between their bright and faint states, while SXTs with accreting neutron stars have a smaller variation. This difference is predicted if the former stars have horizons and the latter have normal surfaces. ", + "introduction": "Soft X-ray transients (SXTs) are binary stellar systems in which a black hole (BH) or neutron star (NS) primary accretes matter from a main sequence or giant secondary (van Paradijs \\& McClintock 1995; Tanaka \\& Shibazaki 1996). A typical SXT displays a large variation in luminosity. Most of the time it remains in a quiescent state and is very dim. In this phase, only a fraction of the mass transferred from the secondary accretes on the primary, the rest being stored in the outer part of an accretion disk. Once every few decades, however, the source goes into outburst and becomes very bright for a few months. During this time, the material stored in the outer disk is apparently accreted very rapidly. In addition to the major outbursts described above, some SXTs also display Type I bursts due to thermonuclear flashes of the accreted material, which unambiguously identifies them as NS SXTs (Joss \\& Rappaport 1984). In a few of these systems the mass of the primary star has been measured and found to be consistent with the mass of a NS ($\\sim1.4M_\\odot$). In several other SXTs, however, the mass of the primary is found to be $>3M_\\odot$, which makes these stars too massive to be NSs. These are identified as BH candidates. Although BHs and NSs have been convincingly distinguished on the basis of their masses, we note that the real physical distinction between the two is that BHs are supposed to have event horizons while NSs are normal stars with surfaces. This basic difference between BHs and NSs has not so far been demonstrated. We show in this paper that SXTs provide a unique opportunity to test the reality of event horizons. Since our argument depends on an understanding of accretion flows, we begin with a discussion of this subject in \\S2. ", + "conclusions": "Figure~2 is a plot of $L_{min}/L_{max}$ {\\it vs.} $L_{max}$ of the BH SXTs and NS SXTs discussed in \\S3. We see a clear confirmation of the basic ideas described in \\S2. First, the BH SXTs all have larger values of $L_{max}$ than the NS SXTs, consistent with their larger masses. Also, none of the NS SXTs has $L_{max}>L_{Edd}=2\\times10^{38} ~{\\rm erg\\,s^{-1}}$ (taking $M_{NS}=1.4M_\\odot$). These facts have been noted before (e.g. Barret et al. 1996). Second, and this is new, without exception every BH SXT in Fig.~2 has a smaller value of $L_{min}/L_{max}$ than every NS SXT. This is exactly what we expect if (1) BHs have horizons and NSs do not, and (2) quiescent SXTs accrete via ADAFs. Note that for some of the BH SXTs we only have upper limits on the quiescent luminosity. When these fluxes are ultimately measured, the difference between NS SXTs and BH SXTs may become even more dramatic. Could the difference in $L_{min}/L_{max}$ merely mean that NS SXTs experience a smaller range of $\\dot m$ between quiescence and outburst compared to BH SXTs? This is unlikely since the two systems are very similar in many respects. Furthermore, NS SXTs are, if at all, likely to have a {\\it larger} swing of $\\dot m$ than BH SXTs. If a NS has a strong enough magnetic field and spins rapidly enough, the field can cause the accreting matter to be flung out through a ``propeller effect'' (Illarionov \\& Sunyaev 1975), thereby dimming the source. Asai et~al. (1996) and Tanaka \\& Shibazaki (1996) have argued for the propeller effect in the NS SXT Cen X--4 in quiescence. If many quiescent NS SXTs undergo the propeller effect, the swing of luminosity between $L_{min}$ and $L_{max}$ will be enhanced in NS SXTs as a class. BHs, on the other hand, cannot have a propeller effect because of the ``no-hair'' theorem which rules out a permanent magnetic field on a BH (Shapiro \\& Teukolsky 1983). Since the presence of the propeller effect in quiescent NS SXTs would tend to wash out the difference in $L_{min}/L_{max}$ which we predict between NS SXTs and BH SXTs, the fact that the predicted difference is seen clearly in the data is especially significant. Could BH SXTs appear less luminous by expelling most of their energy through outflows? This is possible in principle, but is somewhat contrived since we need a mechanism that is extremely sensitive to mass, switching on suddenly for accretors more massive than $1.4M_\\odot$. We suggest that the most natural explanation for the difference in luminosity swing between BH SXTs and NS SXTs is that BHs have event horizons and NSs do not. It is a basic property of a BH event horizon that it will hide any thermal energy which falls through it. A NS, on the other hand, does not have a horizon and must re-radiate whatever thermal energy it accretes. We suggest that Figure 2 confirms this difference. The argument presented here is robust since it makes use of one of the most basic observables in astronomy, namely the total received flux." + }, + "9701/astro-ph9701143_arXiv.txt": { + "abstract": "We report on the first results from a redshift survey of a flux-limited sample of X-ray clusters selected serendipitously from the ROSAT PSPC data archive. We spectroscopically confirm $\\highz$ clusters in the range $0.3$ 20) are associated with extremely compact molecular distributions in the nuclei of merging galaxies. Relative to \\twco, \\thco\\ 1--0 is brightest in quiescent regions of low \\twco\\ surface brightness and weakest in starburst regions and the galactic nuclei. A medium consisting of dense ($n=10^4 - 10^5$ $\\cmmd$) and warm ($T_{\\rm k} > 50$ K) gas will reproduce the extreme line ratios observed in the nucleus of IC~694, where the area filling factor must be at least 20\\%. ", + "introduction": "The inner kiloparsecs of starburst and interacting galaxies harbor stunning amounts of molecular gas, $10^9 - 10^{10} M_{\\odot}$ (e.g., Scoville \\etal\\ 1991; Bryant \\& Scoville 1996). In these environments, molecular clouds are subject to intense radiation fields, supernovae explosions, winds from newborn hot stars, strong tidal forces, and gas surface densities several order of magnitudes higher than in the Milky Way disk. These are also extremely active star formation sites. Knowledge of the physical conditions and structure of molecular gas in interacting systems is essential to understand the starburst activity and its role in galaxy evolution. Arp~299 is an IR-luminous ($L_{\\rm IR} \\approx 8\\times10^{11}$ L$_{\\odot}$) merger system of two galaxies, IC~694 and NGC~3690. Strong \\twco\\ emission has been detected from the nuclei of IC~694 and NGC~3690 and from the interface between the two galaxies (Solomon \\& Sage 1988; Casoli \\etal 1989; Sargent \\etal\\ 1987; Sargent \\& Scoville 1991). The two nuclei, as well as the western overlap region, currently harbor intense star formation activity (c.f., Gehrz \\etal\\ 1983; Baan \\& Haschick 1990). Furthermore, the nucleus of IC~694 is a flat-spectrum radio source, and may be an AGN (Gehrz \\etal\\ 1983). Lower resolution (single dish) observations reveal an unusually large \\twco/\\thco\\ 1--0 line intensity ratio, $\\gapprox 20$ in Arp~299 (Aalto \\etal\\ 1991; Casoli \\etal\\ 1992). These observations left it unclear whether this is due to weak \\thco\\ in the whole system or to a varying \\twco/\\thco\\ 1--0 line ratio. Observations at 20$''$ and 11$''$ resolution by Casoli \\etal (1992) suggest little variation in the \\twco/\\thco\\ 1--0 line ratio and none in the \\twco/\\thco\\ 2--1 line ratio. In contrast, Aalto \\etal (1995) note substantial variations at 28$''$ resolution in the \\twco/\\thco\\ 2--1 ratio: the ratio is about 17 in IC~694; close to 9 in NGC~3690; and about 7 in the interface region between the two disks. Solomon \\etal\\ (1992) detected bright HCN emission in 28$''$ maps of Arp~299. They measured \\twco/HCN ratios of 11 in IC~694 and 13 in the interface region. \\section {Observations and Results} Aperture synthesis \\twco, \\thco, and HCN 1--0 mapping of Arp~299 was carried out with the Owens Valley Radio Observatory (OVRO) millimeter array between March, 1995, and February, 1996. SIS receivers on the six 10.4 m telecopes provided typical system temperatures (SSB) of 600 K, 450 K, and 350 K for \\twco, \\thco, and HCN. Quasars 1150+497 and 0917+449 were used for phase calibration and Uranus and Neptune for absolute flux calibraton. The synthesised beams are $2.''5 \\times 2.''2 $ for \\twco\\ (uniform weighting), $4.''3 \\times 3.''6 $ for \\thco\\ (natural weighting), and $5.''6 \\times 5.''3$ for HCN (natural weighting). At 2.6 mm wavelength with $2.''3$ resolution, a brightness temperature of 1 K corresponds to 57 mJy beam$^{-1}$. The digital correlator, centered at $V_{\\rm LSR} = 3100$ \\kms, provided a total velocity coverage of 1123 \\kms\\ for \\twco, 1175 \\kms\\ for \\thco, and 1407 \\kms\\ for HCN. Data were binned to 4 MHz resolution, corresponding to 10 \\kms\\ for \\twco\\ and \\thco\\ and 13 \\kms\\ for HCN. At 110 GHz an unresolved continuum source of $17 \\pm 2$ mJy was detected in the center of IC~694. Continuum emission was also detected in NGC~3690 ($5 \\pm 2$ mJy) and the overlap region ($9 \\pm 2$ mJy). We subtracted this continuum emission from the line emission before maps were made. The main structures found by Sargent \\& Scoville (1991) with the three telescopes array are recovered in our new \\twco\\ map (Figure 1a), but our increased $uv$ coverage enable improved deconvolution. Lower surface brightness emission (A2), possibly a molecular disk or bar, extends 10-15$''$ southeast of the IC~694 nucleus (A1), coincident with the remnant optical disk. The center of NGC~3690 (B1) is also surrounded by weak, extended emission (B2 and B3) with a somewhat S-shaped morphology. Where the two galaxies overlap, three distinctive clumps (C1, C2, and C3) can now be discerned. Weaker, extended emission is also recovered better in our new \\twco\\ map and clumpy structures can be distinguished at the center of the map (F). These structures appear to connect the major components A, B, and C. In addition, there is a clump (D) north of the main features with systemic velocity $\\approx$ 3280 \\kms, which is beyond the bandwidth of the earlier OVRO data. Bright 6 cm radio continuum peaks at the nuclei of IC~694 and NGC~3690 (Gehrz \\etal\\ 1983; Condon \\etal\\ 1991) coincide with the \\twco\\ peaks to within the estimated positional uncertainty ($0.''5$). There is also reasonable positional agreement between the two brightest \\twco\\ clumps in the overlap region (C) and two additional 6 cm radio continuum peaks: C1 and the western radio continuum peak are also coincident within $0.''5$, although the discrepancy between C2 and the eastern radio continuum peak is somewhat larger, $\\Delta \\alpha = 0.''8$. Weaker radio continuum emission at 18 cm and 6 cm is spatially coincident with the extended \\twco\\ emission in regions A2, B2, B3, and F (Baan \\& Haschick 1990; Gehrz \\etal\\ 1983). The velocity field image (Figure 2, Plate 1) reveals a monotonic shift from the blueshifted emission from A2 in IC~694 to redshifted emission from the overlap region C and from region D. Velocity gradients within C and D are small. The velocity field of NGC~3690 is complicated by a double-peaked emission structure of B2, and both B2 and B3 appear blueshifted relative to the center, B1. The nuclei of both IC~694 (A1) and NGC~3690 (B1) remain unresolved by our $2.''2$ synthesised beam. A two-dimensional Gaussian fit to the nuclear emission of IC~694 and NGC~3690 yields upper limits to the source diameters of $1.''4$ and $1.''5$, corresponding to radii of 140 and 150 pc, respectively (for $D$=42 Mpc). This implies a lower limit of 18 K to the \\twco\\ brightness temperature in the IC~694 nucleus, and therefore, the cloud filling factor is quite high in the inner 200 pc. Even if the intrinsic brightness temperature is as high as 100 K, the surface filling factor of clouds is still almost 20\\%. The three features in the overlap region are resolved, with sizes of $792 \\times 322$ pc (C1), $876 \\times 428$ pc (C2), and $684 \\times 456$ pc (C3). Derived properties for all designated regions are presented in Table 1. The total molecular mass of Arp~299, estimated from $M({\\rm H}_2)=1.2 \\times 10^4 S \\Delta v D^2$ M$_{\\odot}$ ($D$ is the distance in Mpc; $S \\Delta v$ is the integrated \\twco\\ 1--0 line flux in Jy K kms$^{-1}$), is $7.5 \\times 10^9$ M$_{\\odot}$, 87\\% of the value estimated by Solomon and Sage (1988) from their single dish map. This formula corresponds to $N$(H$_2$)/$I$(\\twco)=$3 \\times 10^{20}$ $\\cmmt$ (K \\kms )$^{-1}$ (Sargent \\& Scoville 1991), a standard Galactic \\twco\\ luminosity to H$_2$ mass ratio. The conversion factor may vary, however, across Arp~299, since the line ratios indicate different cloud properties. The most striking feature of the \\thco\\ 1--0 map (Figure 1b) is the {\\it absence of strong emission at the nucleus of IC~694 (A1)}. Emission {\\it is} detected in the A2 region of IC~694, in the overlap region C, and in NGC~3690. Unlike Casoli \\etal\\ (1992) we find significant variation in the \\twco/\\thco\\ 1--0 line ratio across Arp~299. >From the very high value of $60 \\pm 15$ at the IC~694 nucleus (A1) the ratio drops to $10 \\pm 3$ in the A2 region. The global ratio for NGC~3690 is $13 \\pm 2$, but there is an indication that the ratio is higher in the nucleus, B1, than in B2 or B3. The ratio for the east (C3) and west (C1) overlap region is $11 \\pm 3$ and $20 \\pm 4$ respectively. The spatial correlation between the \\thco\\ emission and radio continuum emission peaks is also poor. As for \\twco, the HCN map (Figure 1c) is dominated by a peak at the center of IC~694 (A1), where the \\twco/HCN line ratio is $9 \\pm 1$. Emission is also detected at the nucleus of NGC~3690, where the ratio is $14 \\pm 3$, and in the western (C1) overlap region, where \\twco/HCN= $11 \\pm 3$. There is very little HCN in regions A2 and C3. The HCN and \\thco\\ line emission peaks appear anticorrelated. \\section {The molecular line-ratios} As a system, Arp~299 is not deficient in \\thco. In all regions but the core of IC~694, we observe \\twco/\\thco\\ line ratios typical of star forming regions in other galaxies (e.g., Aalto \\etal\\ 1995, 1991; Young \\& Sanders 1986; Rickard \\& Blitz 1985). Relative to \\twco, \\thco\\ 1--0 is brightest in quiescent regions of low \\twco\\ surface brightness, and weak in starburst regions and galactic nuclei. In contrast, HCN 1--0, like \\twco, is bright in the two galaxy centers and in regions of active star formation. These line ratio variations are, most likely, caused by differences in line excitation. Our results support the suggestion (Aalto \\etal\\ 1995) that {\\it unusually} high (i.e. $>20$) \\twco/\\thco\\ 1--0 line ratios tend to be associated with extremely compact molecular distributions centered on the nuclei of merging galaxies and are primarily due to a small or moderate optical depth, $\\tau \\lapprox 1$, in the \\twco\\ 1--0 line. High ambient pressures, strong tidal forces and ongoing starburst or AGN activity lead to substantial changes in cloud structure and physical conditions. Bryant (1996) also finds that HCN 1--0 is bright in regions of high \\twco\\ surface brightness in merging galaxies. \\subsection{IC~694} The line ratio variation from the IC~694 disk, where \\twco/\\thco\\ = 10 and \\twco/HCN $\\gapprox$ 25, to its nucleus, where the corresponding values are 60 and 9, reflects a dramatic change in cloud properties. The bright HCN line accompanied by relatively weak \\thco\\ emission (HCN/\\thco\\ = $7 \\pm 2$) implies a population of unusually dense and warm clouds. The HCN 1--0 strength implies densities $n \\gapprox 10^4$ $\\cmmd$ if the HCN excitation is dominated by collisions with H$_2$. It is also likely that the density is $\\lapprox 10^5$, so that most of the HCN population will remain in the lower levels and $\\tau_{10} > 1$. At these densities, the \\twco\\ and \\thco\\ 1--0 transitions are thermalised. If the kinetic temperature is also high, the lower levels may become significantly depopulated, effectively reducing the optical depth of the 1--0 line. Then the \\thco\\ 2--1/1--0 line ratio should be $>$ 1. Comparing the single dish \\thco\\ 2--1 flux (Aalto \\etal\\ 1995) with our \\thco\\ 1--0 flux from the nucleus of IC~694, we estimate \\thco\\ 2--1/1--0 $\\gapprox 2$, implying that the gas temperature is high, $>$ 50 K. Although the single dish beam was large, 28$''$, the bulk ($ \\gapprox $65\\%) of the \\twco\\ 2--1 emission within the beam originates in the nucleus of IC~694 (A1). Since the lower transitions of \\twco\\ and \\thco\\ appear thermalised, LTE can be assumed to estimate the \\twco\\ column density, $N$(\\twco). We find it unlikely that the \\twco/\\thco\\ abundance ratio is extremely high (section 3.3). Therefore, a high \\twco/\\thco\\ 1--0 line ratio indicates a low to moderate optical depth ($\\tau \\lapprox 1$) in the \\twco\\ 1--0 line. The high intrinsic \\twco\\ 1--0 brightness temperature makes $\\tau << 1$ unlikely, and we therefore assume $\\tau_{10}(^{12}{\\rm CO}) \\approx 1$. The optical depth of the 1--0 transition can be expressed as $\\tau_{10}(^{12}{\\rm CO}) \\approx 3.9 \\times 10^{-15} N(^{12}{\\rm CO}) (1- e^{-5.53/T_{\\rm ex}})/T_{\\rm ex} \\Delta V$. For a temperature $T_{\\rm ex}$=100 K, line width $\\Delta V = 5$ \\kms, and $\\tau_{10}$=1, the \\twco\\ column density $N(^{12}{\\rm CO}) = 2 \\times 10^{18}$ $\\cmmt$ (per cloud) and the resulting brightness temperature $T_B(^{12}{\\rm CO}) \\approx 60$ K. For a density of $n= 10^4$ $\\cmmd$ and a \\twco\\ abundance, [\\twco/H$_2$]=$5 \\times 10^{-5}$, the cloud radius $r=N(^{12}{\\rm CO})/2x(^{12}{\\rm CO})n({\\rm H}_2)$=0.7 pc --- not unlike cores or clumps within Giant Molecular Clouds in the Galaxy. At this gas density, HCN is not thermalised so $T_{\\rm ex}({\\rm HCN})$ for the 1--0 line will be significantly lower than that for \\twco. If the abundance ratio [\\twco/HCN] $\\approx 10^3$, then $\\tau_{10}$(HCN) is an order of magnitude higher than $\\tau_{10}$(\\twco). The resulting $T_{\\rm B}$(HCN) is sufficiently high to account for a \\twco/HCN line ratio of 9. Above, we infer a clumpy molecular medium because we chose a $\\Delta V$ = 5 \\kms, yielding small clouds with $r < 1$ pc (cf Aalto \\etal 1991). We can not, however, exclude significantly larger $\\Delta V$ which would indicate a continous, non-cloudy structure --- perhaps even a smooth, rotating disk. A third alternative is a molecular ISM consisting of dense clumps surrounded by diffuse, non-cloudy molecular gas (e.g. Aalto \\etal 1994,1995; Dahmen \\etal 1996). \\subsection{The overlap region C and NGC~3690} Molecular line ratio variations are also seen within the overlap region C, albeit smaller than those within IC~694. The weakest \\thco\\ and strongest HCN 1--0 emission, relative to \\twco, is found in C1, the location of the brightest H$\\alpha$ emission in Arp~299 (Gehrz \\etal\\ 1983). Continuum emission at 3.4 $\\mu$m and 10$\\mu$m also peaks close to the C1 clump. Thus, C1 appears to be currently the most active starforming region within C. We suggest that the observed molecular line ratio gradients are the result of a temperature and/or density gradient across the overlap region. The \\twco/\\thco\\ 1--0 line ratio in C1 is considerably lower than in the nucleus of IC~694, perhaps because C1 is an extranuclear starburst. Unlike \\twco, the \\thco\\ emission peak is not connected with the nucleus of NGC~3690 (B1). The radial excitation gradient is similar, therefore, to that of IC~694, with the highest \\twco/\\thco\\ line ratio in the central region. The nucleus of NGC~3690, with intense associated H$\\alpha$ emission, is a site of starburst activity (Gehrz \\etal 1983). \\subsection{Molecular abundances} It has been suggested that high \\twco/\\thco\\ intensity ratios in mergers are caused by unusually high isotopic abundance ratios in molecular clouds with optically thick \\twco\\ 1--0 lines. An influx of very low metallicity gas from the outer disk of the galaxies is the proposed cause of such an extreme abundance ratio (e.g. Casoli \\etal 1992; Henkel \\& Mauersberger 1993). Since the \\twco/\\thco\\ intensity ratio is normal in the outskirts of IC~694 and NGC~3690 this scenario is unlikely to be the explanation for Arp~299. Instead, the measured line ratio variations most likely indicate differences in the line excitation and gas properties in different parts of the system. The observed \\twco/\\thco\\ 1--0 line ratio is, however, a lower limit to the \\twco/\\thco\\ abundance ratio in the emitting region and for $\\tau_{10}$(\\twco) $\\lapprox 1$ this implies an abundance ratio not much greater than 60 in the center of IC~694. This value is typical for GMCs in the Galactic disk, but higher than in the inner region of our Galaxy, where the ratio is $\\approx 25$ (Langer \\& Penzias 1990). Perhaps the ISM in the nucleus of IC~694 recently arrived from the disk of the galaxy. In this case, the difference in line ratio between A2 and A1 is solely caused by a dramatic change in mean optical depth of the \\twco\\ line. On the other hand, selective photodissociation of \\thco\\ by a starburst and/or an AGN may change the isotopic abundance ratio. A young nuclear starburst may also produce extra $^{12}$C and thus temporarily increase the $^{12}$C/$^{13}$C abundance ratio (e.g., Henkel \\& Mauersberger 1993). ", + "conclusions": "The \\twco/\\thco\\ 1--0 line ratio varies dramatically within Arp~299, from $60 \\pm 15$ at the nucleus of IC~694 to 5-10 in its disk and in the eastern and north interface regions (C3 and D). The \\thco\\ 1--0 brightness, relative to \\twco, is high in quiescent regions of low \\twco\\ surface brightness, and low in starburst regions and galactic nuclei. In contrast, HCN 1--0 is bright in the two galaxy centers and in the active extranuclear star formation region. The \\twco/HCN 1--0 is $9 \\pm 1$ at the nucleus of IC~694, $14 \\pm 3$ for NGC~3690 and $11 \\pm 3$ for the extranuclear starburst region C1. Unusually high \\twco/\\thco\\ line ratios ($>$ 20) appear to be associated with extremely compact molecular distributions in the nuclei of merging galaxies (cf. Aalto \\etal\\ 1995). The large \\twco/\\thco\\ 1--0 intensity ratio at the nucleus of IC~694 can be attributed to low to moderate optical depth ($\\tau \\lapprox 1$) in the \\twco\\ 1--0 line, possibly combined with abundance effects. A medium consisting of dense ($n=10^4 - 10^5$ $\\cmmd$), warm ($T_{\\rm k} > $ 50 K) gas is consistent with the observations." + }, + "9701/hep-ph9701420_arXiv.txt": { + "abstract": "Electron polarization induced by magnetic fields can modify the potentials relevant for describing neutrino conversions in media with magnetic fields. The magnitudes of polarization potentials are determined for different conditions. We show that variations of the electron polarization along the neutrino trajectory can induce resonant conversions in the active-sterile neutrino system, but cannot lead to level crossing in the active-active neutrino system. For neutrino flavour conversions the polarisation leads only to a shift of the standard MSW resonance. For polarizations $\\lambda \\lsim 0.04$ the direct modifications of the potential (density) due to the magnetic field pressure are smaller than the modifications due to the polarization effect. We estimate that indeed the typical magnitude of the polarization in the sun or in a supernova are not expected to exceed $10^{-2}$. However even such a small polarization may lead to interesting consequences for supernova physics and for properties of neutrino signals from collapsing stars. ", + "introduction": "\\vskip 0.3cm Neutrino propagation in magnetized media has attracted considerable attention recently, both from the point of view of the early universe cosmology as well as astrophysics. The presence of magnetic fields in the Universe \\cite{Zeldo} as well as in various astrophysical objects \\cite{Raffelt0} can affect neutrino conversion rates and this could have important implications. It has been shown that the neutrino dispersion relations in media with non-zero magnetic fields are modified with respect to those in vacuo and that the effect of the magnetic field can be equivalently described as a correction to the neutrino self-energy \\cite{disp,Raffelt}. An alternative equivalent approach to this problem has been given in ref. \\cite{SemikozValle} where the matrix element of the axial vector current has been calculated for an electron-positron plasma in the presence of a magnetic field. The basic result is that the potential relevant for neutrino propagation acquires a component (axial potential) which is proportional to the scalar product of the neutrino momentum and the magnetic field vector. In particular, it has been claimed that the axial potential can be larger than vector current term, thus inducing the possibility of a new type of resonant neutrino conversions \\cite{ssv,DOlivo}. Several papers have considered the implications of the axial potential in the propagation of neutrinos in media which may be magnetized either by regular or random magnetic fields. In the latter case neutrino conversions become aperiodic \\cite{SemikozValle}. The effect of such axial potentials on active-sterile supernova and early universe neutrino conversions has been discussed \\cite{sergio}. In addition, the corresponding effect of strong random magnetic field upon neutrino transitions induced by a transition magnetic moment in the early universe hot plasma or in a supernova was discussed in ref. \\cite{sergio1}. It has also been realized recently \\cite{Kusenko} that even small polarization effects may have very remarkable implications. For example, they may lead to an explanation for the birth velocities of pulsars as resulting from asymmetries due to neutrino conversions in the cooling protoneutron star. In this paper we consider some general features of the neutrino propagation in polarized media. We generalize the results of an early computation of the polarization effect in ref. \\cite{Langacker}. In particular, we show that the previously considered magnetization effects can be equivalently treated as the effect of {\\it polarization} of the electrons induced by the magnetic field. We find the potentials in terms of the averaged polarization of the medium (sect. 2) and then calculate the averaged polarizations in the magnetic field for various physical conditions (sect. 3.). This approach gives a more transparent physical interpretation of the results and allows one to obtain an important restriction which was missed in ref. \\cite{ssv,DOlivo}, a fact which led to some incorrect statements. In sect. 4 we study the influence of the polarization on neutrino conversions. We show that electron polarization can induce resonant conversions in the active-sterile neutrino system, but can lead only to a shift of the usual MSW resonance \\cite{MS,Wolfenstein} for neutrino flavour conversions. In sect.5 we consider possible implications of the polarization effects for solar and supernova neutrinos. We have explored quantitatively the expected magnitude of the polarization which is consistent with realistic density, $Y_e$ and temperature profiles found in the sun or a supernova. We have estimated that for polarizations $\\lambda \\lsim 0.04$ the modifications in the potential are smaller than the corresponding direct modifications of the potential due to field pressure. We estimate the typical magnitudes of the polarization for various physical situations. Even though we find that the expected values of the polarization is small, it can lead to interesting consequences for supernova physics and for properties of neutrino signals from collapsing stars. ", + "conclusions": "\\vskip 0.3cm \\noindent \\ben \\item We have shown that neutrino propagation in a magnetized medium \\cite{disp,ssv,DOlivo,sergio,sergio1} can be equivalently seen as being associated with the scattering of neutrinos on electrons polarized by the magnetic field. This approach gives a more transparent physical interpretation of the effect and allows one to obtain important restrictions. \\item We have shown that, for neutrino flavour conversions the polarization term of the potential can not overcome usual vector current term. The polarization term can lead to lowering or raising (depending on the direction of polarization) of the the resonant density and, therefore, shift the position of the resonance layer. In contrast, in the case of active-sterile neutrino mixing, the polarization term can be bigger than the vector term. Thus the polarization can induce resonant neutrino conversions, even for small values of the parameter $\\Delta$ in \\eq{initial}. Moreover, in media with fixed chemical composition the resonant conversions can take place both in neutrino and anti-neutrino channels. \\item For realistic magnetic fields in the sun or in a supernova the polarization does not exceed (0.1-1) \\%. Although small, the shift of the resonance layer can lead to an explanation of observed peculiar velocities of pulsars. \\item Strong magnetic fields may lead via pressure to a direct perturbation of density and therefore, the potential profile. The effect of such direct perturbation becomes stronger than the polarization effect we have studied for $\\lambda \\gsim 0.04$. \\een \\newpage \\noindent {\\bf Acknowledgment}\\\\ \\vskip 0.3cm We thank E. Akhmedov and A. Kusenko for discussions. This work has been supported by DGICYT under Grants PB95-1077 and SAB94-0325, by a joint CICYT-INFN grant, and by the TMR network ERBFMRXCT960090 of the European Union. H. N. has been supported by Ministerio de Educaci\\'on y Ciencia. V. S. also got support from RFFR under grant N. 95-02-03724. \\vskip 1truecm" + }, + "9701/astro-ph9701053_arXiv.txt": { + "abstract": " ", + "introduction": "The open inflationary universe scenario is one of exciting current topics among the physics in the early universe. The possibility to create an open universe $(\\Omega_0<1)$ through bubble nucleation in the context of inflationary cosmology has become under discussion rather recently,\\cite{Got,BGT,YST95,Lindea,Lindeb,GreLid} and several models of inflaton potential have been proposed.\\cite{BGT,Lindea,Lindeb,GreLid} Now a central issue is if these models are compatible with the observed anisotropies of cosmic microwave background (CMB) on large angular scales. In several recent papers,\\cite{YTS95,HAMA,STY95,ST96,YST96,YB,Garriga,Bellido,Cohn} quantum fluctuations of the inflaton field that generate the initial curvature perturbations have been evaluated and the resulting spectrum of CMB anisotropies has been calculated. But a drawback of all the previous studies is that the gravitational degrees of freedom have not been taken into account. The incorporation of gravitational perturbations causes two effects. One is the coupling between perturbations of the inflaton field and those of the metric, which may alter the spectrum of the initial curvature perturbations drastically. The other is the contribution of gravitational wave perturbations to the CMB anisotropy, which has not been taken into account at all in the previous analyses. As has been known, a constant time hypersurface in an open inflationary universe is not a Cauchy surface of the whole spacetime.\\cite{STY95} Thus we cannot set commutation relations on this hypersurface when we consider quantization of a field in the open universe. This difficulty has been solved in the case of a scalar field\\cite{STY95,ST96} and recently a method to manage the gravitational wave modes in the Milne universe has been developed by the present authors\\cite{TanSas} (Paper I). In this paper, we extend the result in Paper I to make it applicable to a general open universe. The aim of this paper is to give the spectrum of gravitational wave perturbations in the context of the open inflationary universe scenario. This paper is organized as follows. In section 2 we review the zeroth order approximation to an open universe model, which is based on the $O(4)-$symmetric bubble nucleation\\cite{Col,ColDeL}, and explain the notation we use in the succeeding sections. In section 3 gravitational wave perturbations of the $O(4)-$symmetric bubble is investigated. First we show a similarity between the massless scalar field perturbation and the gravitational wave perturbation. Then using this similarity, we reinterpret the results for massless scalar field perturbations obtained in Ref.~\\citen{ST96} and give the spectrum of gravitational wave perturbations. Section 4 is devoted to summary and discussion. In this paper, we use the units, $c=\\hbar=1$, and adopt the metric signature, $(-,+,+,+)$. \\newpage ", + "conclusions": "{}In this paper we have derived the spectrum of gravitational wave perturbations in the context of the open inflationary universe scenario. We have assumed that the quantum state after bubble nucleation is given by ``the Euclidean vacuum state'' that is determined by the analyticity of modes when they are continued to the Euclidean region. Under this assumption and in the thin wall approximation, we have explicitly obtained the spectrum (\\ref{spec}). An important feature of the spectrum is that it is infrared finite as opposed to the case of pure de Sitter background. We have also found that there is no discrete spectrum that comes from supercurvature modes. At a glance, there seemed to exist supercurvature modes at $p^2=-1$ but it is shown to be an illusion due to gauge degrees of freedom. A subtlety associated with the even parity $p^2=0$ modes that they become degenerate with the $p^2=-4$ scalar perturbation modes has been also resolved. Taking account of all the degrees of freedom of metric perturbations, we have shown that these modes do not exist. In the previous analyses\\cite{HAMA,Garriga,YST96} without taking into account the gravitational degrees of freedom, the scalar $p^2=-4$ modes occupied a special position because they existed independent of the detail of the potential of tunneling scalar field and were called wall fluctuation modes. Our result implies that these modes cease to exist once the gravitational degrees of freedom are taken into account. This result is consistent with that obtained by Kodama et al.\\cite{KodIsh} in the case of infinitely thin domain wall with vanishing potential energy in both vacua. We should note, however, that our result does not exclude the possibility that a discrete mode describing the wall fluctuation exists with a shifted eigenvalue other than $p^2=-4$." + }, + "9701/astro-ph9701029_arXiv.txt": { + "abstract": " ", + "introduction": " ", + "conclusions": "The hope that complementary water and ice detectors will soon probe Northern and Southern hemispheres for cosmic sources of neutrinos is not unrealistic. Depending on the results, high threshold detectors, constructed in the next decade, can subsequently be expanded or, alternatively, back-filled with additional OMs to reach lower thresholds." + }, + "9701/hep-ph9701214_arXiv.txt": { + "abstract": "Core collapse supernovae are dominated by weakly interacting neutrinos. This provides a unique opportunity for macroscopic parity violation. We speculate that parity violation in a strong magnetic field can lead to an asymmetry in the explosion and a recoil of the newly formed neutron star. We estimate the size of this asymmetry from neutrino polarized neutron elastic scattering, polarized electron capture and neutrino-nucleus elastic scattering in a (partially) polarized electron gas. ", + "introduction": "\\label{introduction} Core collapse supernovae emit approximately $10^{58}$ neutrinos. Apart from gravity, these neutrinos only have weak interactions---whose hallmark is parity violation. Can this lead to {\\it macroscopic} parity violation? The important role played by weak interactions makes supernovae (or related phenomena) unique among all macroscopic systems in the present universe. No other systems have this potential for large scale parity violation. One parity violating observable is a correlation ${\\bf k_\\nu}\\cdot {\\bf B}$ between the neutrino flux in the direction ${\\bf k_\\nu}$ and the direction of the magnetic field ${\\bf B}$. As a result, the entire neutron star, formed in the explosion, could recoil. This can be at a significant velocity because of the tremendous momentum of the neutrinos (see below). Indeed, the observed velocities of neutron stars are large: with three dimensional galactic velocities of the order of 500 km/s \\cite{vel}. Perhaps the simplest explanation for such large velocities is that neutron stars receive a significant kick at birth from an asymmetric supernova explosion. Note that this asymmetry need not be solely from parity violation. For example, one could have a preexisting asymmetry in the collapsing stellar core as speculated by Burrows and Hayes~\\cite{burrows}. These authors also discuss how an asymmetry may impact a detectable gravitational radiation signal. In any case, there is a simple relation between the final velocity of a neutron star and the asymmetry of the neutrinos $A_\\nu$ and other matter $A_m$ ejected. Note, a supernova ejects the mantle and outer envelope comprising some 90 percent of the original star's mass. The binding energy per nucleon of a neutron star is about 100 MeV/n. Thus, neutrinos carry off 100 MeV/c of momentum per nucleon of the neutron star (which then has a gravitational mass of the order of 839 MeV/n). In contrast, other matter only carries off about one percent of the binding energy (1 MeV/n). However, the mass of this material is about 10 times larger than that of the neutron star, so that its momentum ($p_{\\rm m}\\approx [2\\cdot 1{\\rm MeV}\\cdot 10\\cdot 939 {\\rm MeV}]^{1/2}\\approx 140$~MeV/c) is comparable to that of the neutrinos.\\footnote{One can ask why do supernovae explode so that the momentum in matter is comparable to that in neutrinos?} The recoil velocity $v$ of a neutron star is approximately, \\begin{equation} v/c \\approx 0.1 ( A_\\nu + A_m), \\end{equation} with $c$ the speed of light and the asymmetry of the neutrinos is \\begin{equation} A_\\nu={|\\sum_i {\\bf p_i}|\\over \\sum_i |{\\bf p_i}|}, \\end{equation} where the sum runs over all neutrinos emitted. There is a similar expression for the asymmetry of the other matter $A_m$. {\\it To reproduce observed velocities of the order of $10^{-3}c$\\ one needs $A_\\nu$---and/or $A_m$---to be about one percent.} Since matter and neutrinos are coupled, one expects both $A_\\nu$ and $A_m$ to be nonzero, i.e., an original asymmetry in one will produce an asymmetry in the other. Thus, the interesting question is which came first (the chicken or the egg) $A_\\nu$ or $A_m$? In this paper we speculate that the original $A_\\nu$ may stem from parity violation in a strong magnetic field. We assume conventional weak interactions for the neutrinos. Others have considered nonstandard neutrino magnetic moments \\cite{magmom} or matter enhanced neutrino oscillations \\cite{osc}. These could enhance the effect. However, one does not require new interactions for a nonzero asymmetry. Thus, our speculation only depends on the existence of a strong magnetic field. A number of Pulsars are thought to have magnetic fields around $10^{12}$ Gauss. However, the dipole field inside a supernova at early times could be stronger, perhaps $10^{14}$ Gauss~\\cite{magnetic}. Finally, it is possible that there are very strong {\\it non-dipole} fields of the order of $10^{16}$ Gauss~\\cite{strongmag}. For example, differential rotation could wind up the ${\\bf B}$ field into a very strong torus configuration. In section II we calculate the neutrino asymmetry that these fields induce because of several parity violating hadronic and electron reactions. We discuss possible enhancements of the asymmetry and conclude in section III. We end this section with a speculative note regarding macroscopic parity violation and the origin of life. Cline suggests~\\cite{cline} that polarized leptons from a nearby supernova could influence homochiral organic molecule formation. This is an interesting but clearly very speculative atempt to explain why Terrestrial life uses almost exclusively L-amino-acid enantiomers (``left handed'' mirror image molecules). Although not directly related, we speculate that supernovae produce another macroscopic parity violating effect. The hypothesis that parity violation leads to an asymmetry and recoil may be more directly observed and tested. ", + "conclusions": "In this section we discuss possible enhancements of the asymmetry and conclude. Spin dependent strong interactions can enhance the magnetic response of neutron rich matter. This could increase the asymmetry from polarized neutrons. For example, Kutschera and Wojcik~\\cite{ferro} speculate that a ferromagnetic state could form because of spin-dependent neutron-proton interactions. Indeed the asymmetry could still be increased by a significant amount, even if a ferromagnetic state did not form. However, this enhancement is unlikely to be an order of magnitude or more unless one is very close to a ferromagnetic transition---which we think unlikely. Therefore, we do not expect very large enhancements from the strong interactions. Alternatively, the nonlinear dynamics in a supernova could enhance small microscopic asymmetries. Indeed, the explosion energy may be very sensitive to the neutrino heating rate. Consider the extreme limit of a heating rate that just fails to produce an explosion. A small asymmetry in the neutrino flux will produce a large asymmetry in the ejected mater by restarting the shock wave on only one side of the supernova.\\footnote{The sensitivity may be somewhat reduced by requiring the shock to reproduce observed (relatively large) explosion energies.} Perhaps a $10^{-3}$\\ asymmetry in the neutrino flux can lead to a $10^{-2}$\\ asymmetry in the ejected matter. This should be explored with simulations incorporating a small microscopic asymmetry. We have considered a number of electron reactions which yield somewhat small asymmetries for perhaps accidental reasons. For example, polarized electron capture produces a reduced asymmetry because of a broad angular distribution. The natural size of an asymmetry is of the order of $eB/k_F^2$ for degenerate conditions or $eB/E_\\nu^2 \\approx 0.6 \\times 10^{-2}\\ B_{14}$\\ for nondegenerate conditions (with $E_\\nu \\approx 10$ MeV the neutrino energy). This is near one percent for $B_{14}$ near unity. Thus, one should search for other electron reactions with large asymmetries. One possibility is $\\nu\\bar \\nu \\rightarrow e^+e^-$\\ which will be examined in future work. Note that Kuznetsov and Mikheev~\\cite{kuz} considered the related $\\nu\\rightarrow \\nu e^+e^-$ in a strong field. They find a relatively small asymmetry in $A_m$\\ of the order of $10^{-5}$\\ $B_{14}$. Perhaps this is because $B$ is necessary to produce, both, the asymmetry and to make the process kinematically allowed. Thus we expect $\\nu \\bar\\nu\\rightarrow e^+e^-$ to be more important because it is allowed even in the absence of $B$. Supernovae, because they are dominated by weakly interacting neutrinos, provide a unique opportunity for macroscopic parity violation. One parity violating observable is a correlation between the neutrino flux and the magnetic field directions. We have examined possible asymmetries in a supernova from known parity violating weak interactions in strong magnetic fields. To explain the large velocities of neutron stars one needs an asymmetry in the radiated neutrinos {\\it and or} the ejected matter of the order of one percent. We have looked at a number of hadronic and electron reactions that give asymmetries of the order of a few times $10^{-4}\\ B_{14}$. These are somewhat small to directly produce the recoil velocity from a dipole magnetic field. However, they only involve known weak interactions and, thus, are clearly present and provide a benchmark to compare with more speculative possibilities. Furthermore, these asymmetries are important if supernovae involve very strong non-dipole magnetic fields (possibly as high as $10^{16}$ Gauss). One should continue to search for new reactions and or enhancements since the natural size of electron reaction asymmetries $eB/E_\\nu^2\\approx 0.6 \\times 10^{-2}\\ B_{14}$\\ could produce a large enough effect." + }, + "9701/astro-ph9701047_arXiv.txt": { + "abstract": "We present optical and near-infrared (NIR) surface brightness and colour profiles, in bands ranging from $U$ to $K$, for the disk and bulge components of a complete sample of 30 nearby S0 to Sbc galaxies with inclinations larger than 50$^{\\rm o}$. We describe in detail the observations and the determination of colour parameters. Calibrated monochromatic and real-colour images are presented, as well as colour index maps. This data set, tailored for the study of the population characteristics of galaxy bulges, provides useful information on the colours of inner disks as well. In related papers, we have used them to quantify colour gradients in bulges, and age differentials between bulge and inner disk. ", + "introduction": "In a recent series of papers we analyze various structural and population parameters of the bulge component of early-to-intermediate type disk galaxies. We did this by obtaining and analyzing accurate surface photometry of a complete sample of inclined spiral galaxies. In Balcells \\& Peletier (\\cite{BP94}, hereafter BP94) we derive optical colour profiles and compare mean colours and colour gradients of bulges to those of ellipticals. We find that bulges have similar or bluer colours than ellipticals of the same luminosity, implying that many bulges contain younger populations, or have lower metallicity, than corresponding ellipticals. In Peletier \\& Balcells (\\cite{PB96a}, hereafter PB96a) we compare the colours of bulges to those of inner disks. These two colours not only scale with each other, a result already shown in BP94; colours of bulges and inner disks are almost identical; the bulge does not appear as a morphological component in colour index maps. Implied age differences are very small, with d$\\log(age)$ ranging from 0 to 0.15. This result, derived from NIR as well as optical data, is not affected by uncertainties due to dust extinction in the disk. Finally, in Andredakis, Peletier \\& Balcells (\\cite{APB95}, hereafter APB95) we develop a 2D bulge-disk decomposition, an extension of Kent's (\\cite{Kent84}) method, and show that surface brightness profiles of bulges follows an $r^{1/n}$ law, with $n$ ranging from 4-6 for S0's to 1 for Sc's. In a future paper (Peletier \\& Balcells, in preparation ) we will discuss optical-infrared colour gradients in bulges and disks; some of these results have already been presented in Peletier \\& Balcells (\\cite{PB96b}, hereafter PB96b). The validity of these results rests first on the quality of the raw data, and second on the attention given to the problem of dust extinction. The latter is described in detail in the previous papers in this series. As for the data, we were fortunate that photometric conditions pervailed during both the optical and NIR observations; flat-fielding, critical for sky subtraction, was accurate to about 0.2\\% of the sky, and a factor 10 better in the near-infrared. The sample is complete for the specified galaxy types and diameter and magnitude ranges. Optical and NIR colours of nearby galaxies contain important clues on the star formation history of spiral galaxies. Not only can one learn about the stellar contents of these objects, but also, by comparing the properties with galaxies at large redshift one can study the evolution of stellar populations. In the last few years several studies of the colours of spiral galaxies have appeared. Terndrup \\etal (\\cite{Terndrup++}) observed 43 galaxies in $J$ and $K$, which was complemented with $r$-band photometry of Kent (\\cite{Kent84}, \\cite{Kent86}, \\cite{Kent87}). De Jong \\& van der Kruit (\\cite{deJ+vdK}) and de Jong (\\cite{deJ}) published $B$, $V$, $R$, $I$ and $K$-band photometry of 86 spiral galaxies. Both samples contain nearby spirals of types S0 - Sc, while de Jong's sample only contains face-on or nearly face-on galaxies. Peletier \\etal (\\cite{P++94}) imaged 37 galaxies in the $K$-band, and combined this with $B$ and $R$-band images of the ESO-LV catalog (Lauberts \\& Valentijn \\cite{ESO-LV}). Their sample only contains galaxies of types Sb-Scd, at somewhat larger distances and smaller, and which are likely to be severely affected by extinction by dust. The current sample is complementary in type to it, and is complementary to de Jong's sample in inclination. The fact that our sample is biased towards early-type spiral galaxies, with large bulges, the high inclination on the sky, the good sampling, the relative novelty of the near-infrared data, and especially the fact that we have good-quality surface photometry in 6 bands, makes this data set ideal to study the colour properties of bulges and inner disks of early-type spiral galaxies. Using the opportunities of a new electronic medium we present here all of the available data in graphical and tabular form. In \\S~\\ref{observations} we describe the sample and the optical and NIR observations. In \\S~\\ref{reduction} we describe the complete process of obtaining optical and NIR colour profiles. In \\S~\\ref{comparisons} we compare our results to NIR aperture photometry and to published NIR colour profiles. A discussion of the data quality is given in \\S~\\ref{discussion}. ", + "conclusions": "\\label{discussion} In this data paper we present high-quality surface photometry for a sample of early-type spiral galaxies in 6 bands, ranging from U (3700 ${\\rm \\AA}$) to K (2.2 $\\mu$m). The galaxies comprise a complete optically selected sample of inclined spirals of type S0 -- Sbc. It is ideal to study the stellar populations of bulges, since the bulges of early-type spirals are large, and in general only one side of the bulges is obscured by the disk, so that the other is relatively free of extinction by dust and of contamination by the disk. The infrared data have been obtained with a 256 $\\times$ 256 InSb array with a field of $\\sim$ 80 $\\times$ 80 arcsec, more than enough to contain the whole bulge, and the inner galaxy disk. The effective seeing is approximately 1$''$ in the infrared and 1.5 $''$ in the optical. Our photometry has been extensively compared with the literature in BP94 (for the optical) and in this paper (for the infrared). For the surface brightness profiles the agreement is good. The absolute calibration is accurate to 0.10 mag in $U$ and $B$, 0.05 in $R$ and $I$, and 0.10 mag in $J$ and $K$. We provide surface brightness and colour profiles for the individual galaxies, as well as measurements of absolute magnitudes, bulge-to-disk ratios, bulge effective radii, disk scale-lengths, mean colours of bulges and disks, colour gradients of bulges and disks, and center shifts in each passbands. The sample may serve as a standard sample for the colours of galactic bulges, replacing the work of Persson \\etal (1979) and Frogel \\etal (1978). It is complementary to the sample of de Jong \\& van der Kruit (\\cite{deJ+vdK}) in inclination, and to the sample of Peletier \\etal (\\cite{P++94}) in galaxy types. It has similar selection criteria as the sample of Terndrup \\etal (\\cite{Terndrup++}). This is the first study of a moderately large number of galactic bulges based on NIR data using 256 $\\times$ 256 arrays. Compared to the recent study of Terndrup \\etal (\\cite{Terndrup++}), based on 58 $\\times$ 62 InSb array data, our spatial sampling in the center is higher, the photometry goes out further, and we cover more passbands (we also give $U$, $B$ and $I$). Colours are treated in more detail in our paper. In particular, we present minor axis colour profiles, which allow for the study of bulge colours unaffected by extinction. We wish to use the electronic format of {\\it New Astronomy} to present all the calibrated images and colour maps electronically for easy access and further processing by the community. We are confident that this will contribute to the furthering of our understanding of the structure and formation history of galaxy bulges." + }, + "9701/astro-ph9701193_arXiv.txt": { + "abstract": " ", + "introduction": " ", + "conclusions": "We have analysed photometry and low-resolution spectroscopy of faint late-K and early M-dwarfs in three intermediate- and high-latitude fields. A comparison with the predictions of a starcount model which matches the faint (R$< 24$) stellar colour-magnitude distribution in two other fields (Reid et al, 1996) shows that these stars, lying at distances of up to 2 kiloparsecs above the Plane and members of the intermediate population. Our analysis, based on the strength of CaH and TiO molecular bands, demonstrates that these stars have abundances close to that of the old disk. Clearly, these results are preliminary, both given the total number of stars in our sample and the approximate nature of the calibration of the measured bandstrength against real abundances. However, our analysis does indicate a possible new direction that can be taken in exploring the abundance distribution within Galactic stellar populations." + }, + "9701/astro-ph9701066_arXiv.txt": { + "abstract": "After $\\sim$16 years of radial-velocity observations of a sample of 22 R-type carbon stars, no evidence for binary motion has been detected in any of them. This is surprising considering that approximately 20\\% of normal late-type giants are spectroscopic binaries, and the fraction is close to 100\\% in barium, CH, and subgiant/dwarf CH and barium stars. It is suggested, therefore, that a process that has caused the mixing of carbon to the surface of these stars cannot act in a wide binary system. Possibly, the R stars were once all binaries, but with separations that would not allow them to evolve completely up the giant and asymptotic giant branchs without coalescing. This coalescence may be the agent which causes carbon produced in the helium-core flash to be mixed outwards to a region where convection zones can bring it to the surface of the star. ", + "introduction": "The Henry Draper classification divided carbon stars into two groups N and R on the basis of their spectral features. The N stars exhibit very strong depression of light in the violet part of the spectrum. They are the classical carbon stars that are most easily detected in infrared surveys and used as tracers of an intermediate age population in extragalactic objects. The R stars, on the other hand, have warmer temperatures, and blue/violet light is accessible to observation, and atmospheric analysis. The most extensive analysis of the R stars has been done by Dominy (1984), who found that in the warm R0$-$R4 stars, C is overabundant on average by approximately 0.7 dex relative to the sun and to normal G and K giants. On the other hand he found that the O abundance is normally near solar, and the N abundance just slightly enhanced. At CNO cycle equilibrium both C and O are depleted substantially. The lack of oxygen and carbon depletion, therefore, led Dominiy to the conclusion that the CNO cycle operating near equilibrium is not responsible for the fact that C/O $>$ 1 in these stars. The excess carbon has to come from some other process. Dominy also found that the s process element abundances are nearly solar, similar to results of Gordon (1968), and Green et al. (1973). This is in sharp contrast with most other carbon and carbon-related stars (see reviews by McClure 1984a, 1985). The N, S, barium, CH, sgCH, and dwarf carbon stars all have significantly enhanced abundances of the s process elements relative to iron (e.g., Lambert 1985; Green and Margon 1994). The N stars and many of the S stars are assumed to have undergone helium-shell flashing on the asymptotic giant branch (AGB), and the third dredge-up phase has brought carbon and s process elements to the surface (e.g., see Iben and Renzini 1983). Like many of the barium, CH, sgCH and dwarf carbon stars, the warmer R stars are too faint to be on the AGB. Scalo (1976) has reviewed the status of various carbon-related stars in the Hertzsprung-Russell diagram, and has pointed out that the R0$-$R4 stars lie in a similar temperature/luminosity domain as the barium stars. They are well below the lower luminosity boundary of the onset of helium shell flashing. The lack of s process enhancement, and low luminosity led Dominy (1984) to conclude that the R stars probably received their enhanced carbon from the helium-core flash at the tip of the first ascent giant branch. This still remains the most viable hypothesis, although it is not known how the core-flash contaminants have been mixed to the surface of the star. ", + "conclusions": "" + }, + "9701/astro-ph9701244_arXiv.txt": { + "abstract": "Moderate-resolution spectroscopic observations from the Keck 10\\m\\ telescope are used to derive internal kinematics for eight faint disk galaxies in the fields flanking the Hubble Deep Field. The spectroscopic data are combined with high--resolution $F814W$ WFPC2 images from the {\\it Hubble Space Telescope} which provide morphologies and scale-lengths, inclinations and orientations. The eight galaxies have redshifts $0.15 \\lesssim z \\lesssim 0.75$, magnitudes $18.6 \\leq I_{814} \\leq 22.1$ and luminosities $-21.8 \\leq M_B \\leq -19.0$ (\\Hconst{75} and $q_0 = 0.05$). Terminal disk velocities are derived from the spatially-resolved velocity profiles by modeling the effects of seeing, slit width, slit misalignment with galaxy major axis, and inclination for each source. These data are combined with the sample of Vogt \\etal\\ (1996) to provide a high-redshift Tully-Fisher relation that spans three magnitudes. This sample was selected primarily by morphology and magnitude, rather than color or spectral features. We find no obvious change in the shape or slope of the relation with respect to the local Tully-Fisher relation. The small offset of $\\lesssim 0.4$ $B$ mag with respect to the local relation is presumably caused by luminosity evolution in the field galaxy population, and does not correlate with galaxy mass. A comparison of disk surface brightness between local and high-redshift samples yields a similar offset, $\\sim$0.6 mag. These results provide further evidence for only a {\\it modest} increase in luminosity with lookback time. ", + "introduction": "Until recently, studies of faint field galaxies have been limited to galaxy counts, colors, and redshift distributions, which can be used to construct luminosity functions at earlier epochs. Such luminosity functions are then compared to models incorporating a certain amount of number evolution (controlled by galaxy formation and merging) coupled with luminosity and color evolution (controlled by star formation histories). Recent models range from those predicting only a small degree of luminosity evolution (\\eg Gronwall \\& Koo 1995) to those invoking entirely new classes of galaxies (\\eg Babul \\& Rees 1992; Babul \\& Ferguson 1996), to those requiring high merger rates (\\eg Broadhurst, Ellis, \\& Glazebrook 1992). A direct measure of luminosity evolution in field galaxies will help to distinguish between various hypotheses. Such measures have recently been attempted, but the results are somewhat contradictory. Schade \\etal\\ (1996a) found evidence for disk brightening by 1.2 $B$ mag in galaxies at redshifts $0.5 \\leq z \\leq 1.2$. Simard \\& Pritchet (1996) found even greater levels of evolution ($2.5 \\pm 0.5$ mag) in a sample of very blue galaxies at $z \\sim 0.35$ for which strong \\fOII\\ lines could be spatially resolved. Rix \\etal\\ (1997), also using kinematic information, derived a brightening of 1.5 mag at $z \\sim 0.25$ for sub-$L^\\star$ galaxies. In contrast, Vogt \\etal\\ (1996) and Bershady (1997), using optical rotation curves, and Forbes \\etal\\ (1996), using line widths, found only small deviations from the local Tully-Fisher (TF) relation for spiral galaxies, implying only modest brightening ($\\sim$0.4 mag) out to $z\\sim 1$. Simard \\& Pritchet postulate that these various results could be reconciled if strong luminosity evolution were present {\\it only} in lower-mass systems. If confirmed, this would be an important factor in understanding the evolution of field galaxies. This Letter introduces well-resolved rotation curves of eight predominantly lower-luminosity galaxies ($L_B \\lesssim L_B^\\star$ = -20.3, see Efstathiou, Ellis, \\& Peterson 1988), selected by morphology as suitable TF candidates. These new observations provide a valuable test of the mass-dependent luminosity evolution hypothesis, particularly in comparison to the higher-mass sample presented by Vogt \\etal\\ (1996; hereafter ``\\vfp''). Combined with the work of \\vfp, these data form a sample of rotation curves for 16 galaxies at redshifts {0.15~$\\lesssim$~\\z~$\\lesssim$~1}, ranging over half the age of the universe (for $q_0 = 0$, one--third for $q_0 = 0.5$). We also use this combined data set to explore trends in surface brightness for comparison with Colless \\etal\\ (1994), Forbes \\etal\\ (1996) and Schade \\etal\\ (1996a). Detailed analysis of a significantly larger data set is currently underway, and these results will be used to explore a variety of relevant selection effects. A full description of our analysis techniques is deferred to that paper (Vogt \\etal\\ 1997). ", + "conclusions": "In summary, we have compared a set of 16 high redshift galaxies with a local TF relation and find a modest amount of luminosity evolution ($\\Delta$M$_B \\lesssim$ 0.4). This conclusion is supported by an examination of the surface brightness characteristics of the sample, which show evidence for evolution at the level of $\\sim$ 0.6 magnitude. We find no evidence for deviation from a linear TF relation at lower luminosities, down to a magnitude M$_B \\sim -$19. The bluest galaxies within the sample have a slightly larger offset from the local TF relation, which suggests (when taken in conjunction with results of other studies) that the derived degree of luminosity evolution may depend strongly on sample selection." + }, + "9701/astro-ph9701072_arXiv.txt": { + "abstract": " ", + "introduction": "The boundary layer region between a star and accretion disc is of fundamental importance for non-magnetic accreting systems. This is because up to half the total accretion energy may be liberated over a relatively small scale in this region. (Lynden-Bell and Pringle 1974, Pringle 1977). Consequently, the angular velocity changes rapidly from a near Keplerian value to a smaller value associated with the accreting star on a scale length that is expected to be comparable to the pressure scale height of the slowly rotating star. In a thin Keplerian disc, the inflow velocity is generally highly subsonic. However, in the boundary layer where the gradients increase the radial infall velocity may become large, reaching supersonic values, if an unmodified viscosity prescription appropriate to the outer disc is used (see Papaloizou and Stanley 1986; Kley 1989, Popham and Narayan 1992). In this case, it has been argued (Pringle 1977) that the star would lose causal connection with the outer parts of the disc so that information about the inner boundary conditions could not be communicated outward. In order to prevent such a situation, the viscosity prescription should be modified so as to prevent unphysical communication of information. Various approaches that limit the viscosity in the vicinity of the star (thus reducing the radial inflow velocity) have been suggested (Papaloizou \\& Stanley 1986, Popham \\& Narayan 1992, Narayan 1992). Here we adopt an approach frequently used in non-equilibrium thermodynamics (eg. Jou, Casa-Vasquez and Lebon 1993), and we assume that the viscous stress components relax towards their equilibrium values on a characteristic relaxation timescale $\\tau.$ This leads naturally to a set of basic equations incorporating a finite propagation speed for viscous information given by $c\\subscr{v} = \\sqrt{\\nu /\\tau},$ where $\\nu$ is the usual kinematic viscosity. We use these to investigate the structure of the boundary layer region between the disc and a comparatively slowly rotating star by studying vertically averaged one dimensional models, as many of their properties are expected to be manifested in the more general two dimensional case. We begin with a study of steady state polytropic disc models and then go on to study time dependent models in which energy dissipation and heat transport are taken into account using, for illustrative purposes, parameters appropriate to protostellar discs. ", + "conclusions": "" + }, + "9701/astro-ph9701187_arXiv.txt": { + "abstract": "We have searched for dense molecular gas in three barred spiral galaxies with young starbursts, NGC~3049, 5430 and 6764, which are known Wolf-Rayet galaxies. We detected HCN in the latter two, and CS was marginally detected in NGC~6764. The dense molecular gas contents of the three galaxies are compared to those of other galaxies and to other indicators of star formation. The HCN luminosities (relative to the CO and far infrared ones) in these galaxies with very young starbursts are consistent with those observed in galaxies with older starbursts and in normal galaxies, and so are our upper limits to the CS intensities (relative to CO). The starburst ages evaluated from our spectrophotometric observations are in the range 3.4 to 6.0 Myr. A circum-nuclear ring is apparent on our images of NGC 5430, the galaxy with the oldest central starburst; this galaxy also has the widest molecular lines. The central star formation rates derived from the H$\\alpha$ luminosity are consistent with those expected from the global FIR luminosities, and are correlated with the HCN luminosities. Finally, an independent estimate of the H$_2$ column density is obtained by optical spectrophotometry; it leads to a H$_2$ column density to CO intensity ratio which is about 2 to 3 times lower than the standard value, because the CO intensities of the three galaxies are higher than average, relative to their far infrared fluxes. ", + "introduction": "\\begin{table*} \\caption[~\\TPARAM]{Basic data and global parameters of the three galaxies. D is the distance, $B_{\\rm tc}$ is the apparent corrected blue magnitude and $i$ is the inclination. The logarithms of total HI mass, and of far infrared and blue total luminosities are given in the last three columns} \\input t1.tex \\end{table*} Despite numerous studies of starburst galaxies, many questions remain about the starburst phenomenon, its triggering and evolution. Gravitational interactions and mergers seem to play a major role in triggering starbursts, but violently interacting galaxies are not necessarily the seat of starbursts (Bushouse 1986) and most starbursts seem to be isolated (Contini 1996, Coziol et al. 1996b). Numerical simulations have shown that the bar plays a major role in this process, by efficiently funneling molecular clouds toward the inner few kiloparsecs of galaxies (Nogushi 1988, Friedli \\& Benz 1993). This is confirmed by radio continuum observations (e.g. Puxley et al. 1988). However, the link between bar and far infrared luminosity (which traces young massive stars) in starburst galaxies remains controversial (Hawarden et al. 1996). Molecular clouds obviously play a crucial role in the process of star formation, and their properties have been extensively studied, via millimeter observations of the CO molecule. A strong far infrared (FIR) luminosity has generally been a successful criterion for detection of the molecule in external galaxies, confirming the link between the two indicators of star formation. A CO--$N$(H$_2$) conversion factor has been proposed (Strong et al. 1988) and validated by observations of normal as well as starburst galaxies (e.g. Sage et al. 1990); further studies have shown that it depends on metallicity (Wilson 1995, Arimoto et al. 1996). But its validity has recently been questioned (e.g. Nakai \\& Kuno 1995). It has recently been emphasized that CO only traces low-density molecular gas and several searches for dense (n(H$_{2}$)~$>10^4$~cm$^{-3}$) molecular clouds in normal and starburst galaxies, generally selected to be strong CO emitters, via detection of HCN, CS, HCO$^+$, and other molecules, have been initiated (e.g. Mauersberger et al. 1989, Nguyen-Q-Rieu et al. 1992, Helfer \\& Blitz 1993, Aalto et al. 1995). It turns out that bulges of normal as well as starburst galaxies contain large quantities of dense gas, and that a threshold in the surface density of dense gas does not seem to be required for violent star formation (Helfer \\& Blitz 1993). In view of these mixed results, we have adopted a global view on starbursts in galaxies, based on multi-wavelength observations of a large and homogeneous sample of barred starburst galaxies. Such an approach should enable us to establish quantitative relationships between the properties of starbursts (age, star formation rate and initial mass function), the neutral (atomic and molecular) gas content and the morphology of the host galaxies. It has in fact already given rise to new and original results (Contini et al. 1995, Contini 1996, Contini et al. 1996a, Coziol et al. 1996a). In this paper, we combine millimetric observations of the dense molecular gas with optical images and long-slit spectroscopy of a few examples of very young starbursts galaxies, namely Wolf-Rayet galaxies, in order to investigate the properties of the dense gas of these galaxies in relation to their optical properties. The images provide morphological and photometric informations on the central regions and the bar, and the spectra are used to determine the starburst ages and star formation rates as well as an independent estimate of the column density of H$_2$. A preliminary report on this research project has been published by Contini et al. (1996b). ", + "conclusions": "We have searched for dense molecular gas in three galaxies, NGC 3049, 5430 and 6764, which share the following common properties. They are the sites of very young ($\\simeq$~5 Myr) starbursts, as evidenced by the presence of Wolf-Rayet stars and by our age determinations, they have a large far-infrared flux, they are barred and highly inclined. HCN has been detected in NGC~5430 and 6764, CS in the latter galaxy only. The HCN/FIR luminosity ratios of the two galaxies appear to be normal, confirming the trend noticed by Helfer \\& Blitz (1993) for normal and starburst galaxies. The upper limit to that ratio for NGC 3049 is also consistent with that trend. The measured intensities of CS or upper limits are also normal relative to that of CO and HCN in the same galaxies, when compared to other surveys of CS in external galaxies. If large amounts of dense molecular gas are required for star formation in burst mode, they no longer exist 5 Myr after the burst has started, presumably because they have been used up to form stars and/or ionized by the intense radiation emanating from massive hot stars. We also find that the HCN/CO integrated intensity ratios are rather low relative to the mean value found in other surveys of HCN in external galaxies. This is due to the fact that our three galaxies have an unusually high CO integrated intensity relative to their FIR flux, compared to other galaxies, {\\it including ones with young starbursts}. In other words, this higher than usual CO intensity is not a general property of galaxies with young starbursts. The standard CO--$N$(H$_2$) conversion factor overestimates the amount of molecular hydrogen by a factor 2 or 3 in our three Wolf-Rayet galaxies. But again, our galaxies have unusual CO intensities. This standard factor is thus probably valid for starburst galaxies in a statistical way, with a large uncertainty for individual estimates. Can the optical properties of our galaxies provide an explanation for their molecular line properties? The {\\it central} star formation rates in the three galaxies, estimated by the luminosity of the H$\\alpha$ line, are not unusual among starburst galaxies, because they are all proportional (with the right factor) to the {\\it global} star formation rates estimated by the FIR luminosities. The FIR luminosity and that of H$\\alpha$ in the center are highest in NGC 5430 and lowest in NGC 3049, reflecting the pecking order for the luminosities in the molecular lines. This is another confirmation that the latter are not unusual in very young starbursts. We note that NGC 5430, the galaxy with the oldest starburst, also has a circum-nuclear ring. Its presence is consistent with the fact that most of the bar is in differential rotation (Sect. 4.1), as such rings form where the rotation becomes differential (Lesch et al. 1990). The timescale of ring formation (a few 10$^8$ yr) is much larger than the age of the central starburst (6.0 Myr) of NGC 5430, and of the young star clusters (some less than 10 Myr) recently detected in other circum-nuclear rings by the {\\it Hubble Space Telescope} (Maoz et al. 1996). The absence of such a structure in NGC 3049 and 6764 should thus not be attributed to the fact that their central starbursts are younger, but to dynamical properties, as their bars appear to be rigidly rotating. One interesting property of NGC 5430 revealed in the present paper which deserves to be explored further is the fact that the CO and HI velocity profiles are asymmetric in opposite ways; there is relatively less HI where there is more CO. One possible reason for this asymmetry is a more efficient conversion of HI to H$_2$ in the Wolf-Rayet region which may explain the existence of the young starburst at that end of the bar. A detailed comparison of the relative distributions of HI and CO in starburst galaxies would certainly lead to a better understanding of the transformation of the gas before and during starbursts. Finally, we point out that the linewidths of the CO lines are correlated to the central starburst ages of the three galaxies. This is a general property of young starbursts, which has been discovered by us (Contini et al. 1996a)." + }, + "9701/astro-ph9701144_arXiv.txt": { + "abstract": "The Sombrero galaxy, NGC 4594, contains the most numerous globular cluster system of any nearby spiral. It is an ideal candidate in which to study the globular clusters and contrast them with those in Local Group spirals. Here we present B and I imaging from the CTIO Schmidt telescope which gives a field-of-view of 31$^{'}$ $\\times$ 31$^{'}$. Using DAOPHOT we have detected over 400 globular clusters and derived their magnitudes, B--I colors and photometric metallicities. We have attempted to separate our sample into disk and bulge/halo globular cluster populations, based on location in the galaxy. There is some evidence that the disk population is more metal--rich than the bulge/halo globular clusters, however contamination, dust reddening and small number statistics makes this result very tentative. We find that the median metallicity of the bulge/halo globular clusters is [Fe/H] = --0.8. This metallicity is consistent with previous estimates based on smaller samples. It is also similar to the metallicity predicted by the globular cluster metallicity -- galaxy luminosity relation. As with our Galaxy, there is no radial metallicity gradient in the halo globular clusters. This suggests that the spheriodal component of NGC 4594 did not form by a dissipational process. ", + "introduction": "Measurement of the abundances of Milky Way globular clusters (GCs) have provided important clues on the formation and chemical enrichment history of the Galaxy. Combined with kinematics, a clear dichotomy of metal--rich disk GCs and metal--poor halo GCs have been identified (e.g. Zinn 1985). The disk and halo GCs have a mean metallicity of [Fe/H] = --0.6 and --1.6 respectively (e.g. Ashman \\& Bird 1993). Globular clusters in our sister galaxy, M31, have been relatively well--studied so that spectroscopic metallicity measurements are available for about half of them (Huchra, Brodie \\& Kent 1991). A bimodal metallicity distribution is not as obvious for M31 but the analysis of Ashman \\& Bird (1993) indicates peaks at [Fe/H] = --0.6 and --1.5, i.e. similar to the Milky Way. The other Local Group spiral, M33, has about 30 known GCs, of which spectroscopic metallicities have been obtained for 22 (Brodie \\& Huchra 1991). The mean metallicity of [Fe/H] = --1.55 $\\pm$ 0.37 is similar to that of Milky Way and M31 halo GCs. The small number of GCs makes it difficult to convincingly identify a population of disk GCs. Beyond the Local Group the nearest spirals are M81 (m--M = 27.8) and NGC 4594 (m--M = 29.73). Perelmuter \\& Racine (1995) estimated from imaging of M81 that it contains $\\sim$ 200 GCs. Spectra for 30 GCs have been obtained (Brodie \\& Huchra 1991; Perelmuter, Brodie \\& Huchra 1995) which give a mean [Fe/H] = --1.48 $\\pm$ 0.19. The Sombrero galaxy, NGC 4594, is a giant Sa galaxy (M$_V$ = --22.0) with almost 2,000 GCs (Bridges \\& Hanes 1992). This is significantly more than either the Milky Way with N$_{GC}$ = 160 $\\pm$ 20 or M31 with N$_{GC}$ = 350 $\\pm$ 100 (Harris 1996). It is also nearly edge--on ($i = 84^{\\circ}$) so that GCs projected close to the galaxy major axis can be used to define a plausible sample of disk GCs. Separating disk and halo GCs in M31 ($i = 38^{\\circ}$) has been a major source of uncertainty. Thus NGC 4594 is an ideal candidate in which to study the GC system of another spiral galaxy. The distance to NGC 4594, from the average of the surface brightness fluctuation and planetary nebulae luminosity function methods, is 8.8 $\\pm$ 0.4 Mpc (Ciardullo, Jacoby \\& Tonry 1993). At this distance GCs range in magnitude from B $\\sim$ 18--27, so that only the brightest few GCs are accessible spectroscopically by today's large telescopes. A valiant effort was made recently by Bridges \\etal (1996). After 7 hours of integration time over 4 nights of observing with the 4.2m William Herschel Telescope, they obtained spectra of 34 confirmed GCs. Unfortunately their spectra were not of sufficient quality to determine a metallicity for individual GCs. Summed together they derived a mean of [Fe/H] = --0.7 $\\pm$ 0.3. An alternative approach, is to obtain GC colors from imaging studies. Imaging has the advantage of being very efficient and photometric metallicities can be derived from the Galactic GC color--metallicity relation (Couture \\etal 1990). The GC system of NGC 4594 has been studied photometrically by Harris \\etal (1984) and Bridges \\& Hanes (1992). The first of these was a wide field-of-view photographic study in U and V bands. They examined the spatial distribution and estimated the GC specific frequency. The second study targeted three areas of the galaxy and imaged these in B and V bands with a 2.2$^{'}$ $\\times$ 3.6$^{'}$ CCD. From the B--V colors of 131 GCs they derived a mean [Fe/H] = --0.81 $\\pm$ 0.25 for the GC system. Here we present 31$^{'}$ $\\times$ 31$^{'}$ field-of-view CCD imaging of NGC 4594 in B and I bands. We measure B--I colors which are about twice as sensitive to metallicity as B--V colors (e.g. Geisler, Lee \\& Kim 1996). Our wide field-of-view imaging allows us to define spatially samples of `disk' and `bulge/halo' GCs, and examine their metallicity distributions. ", + "conclusions": "In this section, we first discuss the GC sample of Bridges \\etal (1996). Second, we examine the GC metallicity distribution in NGC 4594 and finally we search for spatial abundance gradients in this distribution. After superposing the position of the Bridges \\etal 34 spectroscopically--confirmed GCs on our CCD images we managed to measure B and I magnitudes for 24 (the others are presumably fainter than our limiting magnitude). The mean color is B--I = 1.9 $\\pm$ 0.2 (error on the mean), which corresponds to [Fe/H] = --0.7 $\\pm$ 0.5. This is reassuringly consistent with the spectroscopic mean metallicity, for all 34 GCs, of [Fe/H] = --0.7 $\\pm$ 0.3. This suggests that the transformation from color to metallicity is reasonable. In Fig. 1 we show the metallicity distribution for disk ($\\vert$Z$\\vert$ $\\le$ 4 kpc) and bulge/halo ($\\vert$Z$\\vert$ $>$ 4 kpc) objects. The median metallicity of the bulge sample is [Fe/H] = --0.8. This is similar to that obtained by Bridges \\etal (1996) and suggests that our sample is dominated by {\\it bona fide} GCs. As pointed out by Bridges \\etal the GC metallicity is more comparable to that found in giant ellipticals than other spirals. However this might be expected given the claimed correlation of GC metallicity with parent galaxy luminosity (e.g. Brodie \\& Huchra 1991). Using the recent best--fit of Forbes \\etal (1996) and the bulge luminosity of NGC 4594, the relation predicts a GC system mean metallicity of [Fe/H] = --0.6 $\\pm$ 0.3 (rms error). The disk sample has a median metallicity of [Fe/H] = --0.5. As the disk sample will include a significant fraction of bulge GCs, the average metallicity of the disk sample is likely to be somewhat higher than --0.5. If we separate the 24 Bridges \\etal (1996) GCs, with available B--I colors, into bulge and disk systems, then their metallicities are [Fe/H] = --1.2 and --0.1 respectively. Thus there is some tentative evidence, from both our sample and that of Bridges \\etal that NGC 4594 contains a population of disk GCs that are slightly more metal--rich than non--disk GCs. Although it is difficult to rule out preferential reddening of the GCs by dust, we do not think this has had a strong effect on the disk GC colors. First, we have excluded the inner galaxy regions with the obvious dust lane. Second, we do not find a strong radial GC color gradient in either scale height (see Fig. 2) or along the major axis as might be expected from centrally concentrated dust. The possibility that inner GCs are instrisically blue (and so removing a reddening trend due to dust) is very unlikely given studies of other galaxies (e.g. Harris 1991; Geisler \\etal 1996). Abundance gradients, or lack of, provide useful constraints to galaxy formation models. For example, the absence of an abundance gradient in Milky Way halo GCs supports the view of Searle \\& Zinn (1978) in which the halo is made up of small protogalactic fragments, rather than arising from a pure dissipational collapse. In the M31 GCs the situation is more uncertain. Huchra, Brodie \\& Kent (1991) found that there exists a upper envelope in GC metallicity with galactocentric radius but also pointed out that selection effects and contamination by disk GCs made evidence for an abundance gradient inconclusive. In Fig. 2 we show the spatial variation of GC metallicity for the disk and bulge/halo samples. In the lower panel of Fig. 2 we show GC metallicity versus projected galactocentric radius for the bulge/halo objects. As with the Milky Way GC system, there is no statistically significant metallicity gradient. This constrains the role of dissipation in forming the spheroidal component of NGC 4594. In a dissipational collapse one expects a metallicity gradient in the direction of the collapse e.g., perpendicular to a disk. Accretion of material may reduce the strength of such gradients. Our Galaxy shows a weak trend of GC metallicity with scale height, where the slope is --0.084 $\\pm$ 0.031 dex kpc$^{-1}$ (Armandroff 1993). In the upper panel of Fig. 2 we show the metallicity of disk GCs versus scale height from the major axis. The data show a hint of a weak trend with scale height, but it is not formally significant." + }, + "9701/astro-ph9701002_arXiv.txt": { + "abstract": "We show that external shocks cannot produce a variable GRB, unless they are produced by an extremely narrow jets (angular opening of $\\sles 10^{-4}$) or if only a small fraction of the shell emits the radiation and the process is very inefficient. Internal shocks can produce the observed complex temporal structure provided that the source itself is variable. In this case, the observed temporal structure reflects the activity of the ``inner engine'' that drives the bursts. This sets direct constraints on it. ", + "introduction": "Five years of BATSE's observations with perfect isotropy and paucity of weak bursts shows that the origin of GRBs is probably cosmological. Therefore, given the measured flux, GRBs involve immense amount of energy $% \\sim 10^{51}$ergs. The ``compactness problem'' then shows that the observed $% \\gamma $-rays must be emitted by a medium with highly relativistic velocities having Lorentz factor $\\gamma \\ge 100$ (Fenimore, Epstein \\& Ho 1993, Woods \\& Loeb 1995, Piran 1995). While the energy source varies from one model to another (binary neutron stars merge, failed supernovae or collapse of magnetic stars) and is relatively speculative, all models of cosmological GRBs involve a relativistic moving shell which converts its (kinetic or magnetic) energy to radiation at a large radius. In all these models the observed radiation does not emerge directly from the ``inner engine'' that drives the shell, which remains hidden. Most bursts are highly variable with a variability scale significantly smaller than the overall duration. Following Fenimore, Madras and Nayakshin (1996), we use kinematic considerations to constrain different GRB models. We show that the overall duration, $T$, reflects directly the length of time that the ``inner engine'' operates and the observed temporal variability reflects variability in the ``inner engine''. The only exceptions to this conclusion are if the engine produces an extremely narrow jet or if GRBs are extremely inefficient. These considerations also limit the emission radius-the place where the energy of the shell is converted to radiation, $R_e$, to be significantly smaller than what was previously thought. The maximal emission radius is quite close to the minimal radius at which a GRB can be produced without becoming optically thick. This is also the place where ``internal shocks'' would naturally take place. Thus, our conclusions are consistent with the ``internal shock'' scenario. Sufficiently small radii are impossible in the hydrodynamic version of the ``external shock'' scenario (and probably in other versions of this scenario as well) . In section 2 we discuss the angular spreading problem, which is the key to our discussion. We show in section 3 that in the framework of models in which the duration of the burst is given by the radius of emission, all solutions to the angular spreading problem result in extremely narrow jets or an extremely low efficiency. In section 4 we discuss models in which the total duration of the burst corresponds directly to the time that the ``inner engine'' operates. The internal shock scenario fits this picture. We show that the hydrodynamic version of the external shock scenario (and most likely all other versions) is incompatible with these limits. ", + "conclusions": "Relativistic motion is essential in GRBs. Relativistic Kinematic arguments (Fenimore, Madras and Nayakshin 1996) strongly limit GRBs models. The duration of a GRB is determined either by the width of the emitting shell, $T=\\Delta /c$, or by the emission radius, $T_{angular}=R_e/\\gamma _e^2$. Models in which $% T_{angular}>\\Delta /c$ (type I) can produce only a single hump smooth bursts. They cannot produce a variable burst. The standard ``external shock'' model is the classical example of a type I model. Therefore, this conclusion rules out this scenario. An exception to this conclusions is if the ``inner engine'' emits an extremely narrow jet (with angular width smaller than $1/N\\gamma _e\\sim 10^{-4}$). Such a jet cannot be produced by a standard fireball in which the matter is accelerated by thermal pressure and it requires another acceleration mechanism. Alternatively a variable burst can be produce by irregular shell with angular fluctuations of large amplitude and small size,or with highly irregular ISM. In this case the process is extremely inefficient due to low covering factor and the total energy needed for a GRB is N times larger than the observed energy (of the order of $10^{53}$ergs). Type II models does not suffer the angular spreading problem and can produce the observed temporal structure. The ``internal shock'' model is the classical example for this type. Note that it is impossible to change the parameters of an ``external shock'' model so that it will become of type II. In this case the temporal structure reflects the activity of the ``inner engine''. The over all duration is the time it operates while the variability reflects the variability of the source. This is good news since we have now a direct information on the ``inner engine''. It is also bad news since only a few known models can produce the observed highly variable temporal structure observed in GRBs." + }, + "9701/astro-ph9701234_arXiv.txt": { + "abstract": "We report on a ROSAT-HRI observation of the soft X--ray transient Cen X-4 during quiescence. We discover a variation in the flux by a factor of 3 in less than four days. This relatively fast variation, the first observed from a quiescent soft X--ray transient, rules out some of the emission mechanisms that have been proposed for the quiescent flux. Accretion either onto the neutron star surface or onto the magnetospheric boundary is clearly favored. ", + "introduction": "The quiescent emission of Soft X--ray Transients (SXRTs) has been investigated for only 5--6 sources (e.g., Verbunt et al. 1994; Asai et al. 1996); this is characterized in all cases by very soft spectra (black body temperatures of a few hundreds of eV) and X--ray luminosities of $\\sim 10^{32-33}$\\ergs. Cen X-4 is one of the nearest ($d\\sim 1.2$ kpc) and best studied SXRT. Two outbursts have been observed from Cen~X-4 in 1969 and 1979. The most detailed observation of Cen X-4 during quiescence is the one recently obtained with ASCA by Asai et al. (1996). These authors reported a quiescent luminosity of $2.4\\times 10^{32}$\\ergs\\ (0.5--10 keV). The X--ray spectrum measured is well fit by a black body component ($T_{\\rm bb}=0.2$ keV; $R_{\\rm bb}=1.8\\times 10^6$ cm) plus a power-law with a photon index of 2--3. The contribution of these two spectral components in the ASCA energy band is comparable. Here we report on a ROSAT HRI observation of Cen X-4 during the quiescent period. ", + "conclusions": "In a reanalysis of the Einstein IPC (on 1980 July 28) and EXOSAT LE (on 1986 February 21) observations of Cen X-4, we find that both are consistent with a single value of the X--ray luminosity for spectral parameters comparable to those of ASCA (Asai et al. 1996), hinting to a constant X--ray luminosity over 15 years. Therefore the decrease in the X--ray luminosity found by van Paradijs et al. (1987) assuming a 1 to 5 keV bremsstrahlung spectrum seems less likely. Our ROSAT HRI observation of Cen X-4 provides also the first evidence for a flux variation on a timescale of a few days in the quiescent flux of a SXRT. This variation is likely intrinsic to the source (no occultations or dips have ever been observed) and therefore it can help constraining the emission mechanisms during quiescence (Stella et al. 1994). The possibility that the late type companion star accounts for the observed luminosity has been excluded by Verbunt (1996). Neutron star cooling is attractive in view of the spectral results consistent with thermal emission from a body of size comparable with a neutron star, however it is incompatible with the observed rapid decrease in the source intensity. Also the shock emission model in which the X--ray luminosity is produced by the interaction of the relativistic pulsar wind with gaseous material from the companion star encounters problems. The mechanism is likely at work in PSR 1259--63, where the X--ray spectrum is characterized by a power-law with photon index of $\\sim 2$ extending from 0.1 to 200 keV (e.g. Tavani \\& Arons 1996). In the case of Cen X-4 however, besides the power-law component, also a black body component is present. Moreover, the HRI count rate variation is mainly observed in the soft channels, where the black body component dominates. We believe that the most likely explanation for the quiescent flux of Cen X-4 is that it derives from mass accretion. In this case variations in the accretion flow can easily explain the observed changes. Either accretion onto the neutron star surface and onto the magnetosphere can explain the observed properties, leaving open the possibility that Cen X-4 contains a fast spinning, weakly magnetized neutron star (Stella et al. 1994)." + }, + "9701/astro-ph9701078_arXiv.txt": { + "abstract": "We have fitted the surface brightness of a sample of 79 elliptical galaxies pertaining to the Coma cluster of galaxies using the \\Se\\ profile. This model is defined through three primary parameters: scale length ($a$), intensity ($\\Sigma_{0}$), and a shape parameter ($\\nu$); physical and astrophysical quantities may be computed from these parameters. We show that correlations are stronger among primary parameters than the classical astrophysical ones. In particular, the galaxies follow a high correlation in $\\nu$ and $a$ parameters. We show that the $\\nu$ and $a$ correlation satisfies a constant specific entropy condition. We propose to use this entropy relation as distance indicator for clusters. ", + "introduction": "The photometrical properties of elliptical galaxies have already been extensively analysed, with special attention paid to the entire range of sizes and luminosities, from giant to dwarf ellipticals. Profile laws are in general defined by various parameters, for example, scale length ($a$), intensity ($\\Sigma_{0}$), and one or more parameters specifying the shape (structure parameters). These parameters can be categorized as {\\em basic} or {\\em primary}, in the sense that they come out naturally from a mathematical definition. The \\DV\\ profile has no structure parameter; for a long time it appeared to be in very good agreement with observations, although hints of systematic deviations were observed; those deviations appear similar for galaxies with the same luminosities (Michard 1985, Schombert J.M. 1986). New observations, especially high resolution data from the HST (e.g. Ferrarese et al. \\shortcite{Ferrarese}), have definitively shown that the \\DV' profile -- rather too inflexible -- does not allow a convenient description of the actual structure of all elliptical galaxies. Some attempts have been made to obtain more general surface brightness laws. Among many possibilities, S\\'ersic's recipe \\shortcite{Sersic}, \\begin{equation} \\Sigma(r) = \\Sigma_{0} \\exp(-u^{\\nu})\\;,\\, u=r/a\\, , \\label{DVlike} \\end{equation} a generalization of the \\DV-law, seems to be better suited to describe elliptical brightness profiles (e.g. Ciotti \\shortcite{Ciotti}, Caon et al. \\shortcite{Caon}, Graham et al. \\shortcite{Graham}, Courteau et al. \\shortcite{Courteau}). In this letter we show the results of the photometrical analysis of a sample of elliptical galaxies in the Coma cluster using a \\Se-model (or a $\\nu$ model). From the basic parameters [$a$, $\\Sigma_{0}$, $\\nu$], it is possible to build physical and astrophysical quantities (effective radius, effective luminosity, total magnitude, energy, entropy, etc\\ldots). All these quantities are in fact a combination of primary parameters; we therefore propose to use primary parameters instead of the classical `astrophysical' quantities in order to obtain better defined correlations. Any correlation between a dimensionless parameter and a distance-dependent one has traditionally led to the definition of a distance indicator; for instance, the Cepheid luminosity--period relation (Luminosity/Period), the Tully-Fisher relation (Luminosity/Rotation), and the Faber-Jackson relation (Luminosity/Velocity dispersion) are widely used. Recently, Young \\& Currie \\shortcite{Young} have pointed out a correlation between the (distance-dependent) intrinsic luminosity and the (dimensionless) shape parameter of the \\Se-law for dwarf ellipticals. A theoretical understanding of the physical underlying processes is crucial to assure the validity of a correlation between the parameters describing an astronomical object. Here, we show that elliptical galaxies are located in the locus defined by constant specific entropy $s$, in the primary parameters plane [$\\nu$, $a$]. The fact that the elliptical galaxies in Comamainly have the same specific entropy should somewhat reflect the intrinsic properties of self--gravitating systems formed through violent relaxation processes. We will discuss whether the entropy relation connecting $\\nu$ and $a$ may be used as a distance indicator. ", + "conclusions": "We have given a physical interpretation for the observed relation between two of the {\\em primary} parameters of the $\\nu$-profile: the values of the $a$ and $\\nu$ parameters of Coma cluster galaxies approximately have a constant specific entropy. The dispersion around the best fit is small, therefore allowing the use of this relation as a distance indicator (relative to the Coma cluster). We must stress that both the observational and theoretical evidences found in this work have to be verified. On the one hand, it is necessary to quantify the universality (i.e. in other clusters) of the ($a$, $\\nu$) correlation found for the galaxies in the Coma cluster. On the other hand, the validity of the assumptions made to derive the theoretical entropy relationship (constant $ M/L$, ideal gas entropy) must be confirmed." + }, + "9701/astro-ph9701016_arXiv.txt": { + "abstract": "We have compared the UV line ratios of a sample of very high redshift radio galaxies (HZRG, $z > 1.7$) with shock and active galactic nuclei (AGN) photoioni\\-zation models, with the goal of determining the balance between jet-induced shocks and AGN illumination in the extended emission line regions (EELR). We find that the UV line ratios cannot be explained in terms of photoioni\\-zation of solar abundance gas by the classical power law of index $\\alpha=$-1.5, which successfully reproduces the general trends defined by the optical line ratios of low redshift radio galaxies. Pure shock models also provide a poor fit to the data. However, photoionization by a power law of index -1.0 provides an excellent fit to the UV line ratios. This suggests that the ionizing continuum spectral shape may depend on radio luminosity and/or redshift, such that it becomes harder as the radio power and/or redshift increase. However, an alternative possibility is that we are seeing the first signs of chemical evolution in these objects, since a power-law of index -1.5 with low metallicity also provides a very good fit to the data. For the high ionization conditions found in the the HZRG, we show that the power-law photoionization mo\\-dels provide a better fit to the data than the shock models. However, such is the complexity of the shock models that we cannot rule out the possibility that a different combination of input parameters can reproduce the observed spectra. We further show that the UV line ratios provide a sensitive test of the ionization mechanism for the lower ioni\\-zation conditions prevalent in some low redshift jet-cloud interaction candidates. For high ionization parameter this discrimination is difficult due to the overlap of shock and power law photoionization models. ", + "introduction": "The character of the emission line spectra of powerful radio galaxies is strongly determined by the ionization mechanism of the gas. For most low redshift ($z<$0.1) radio galaxies photoionization by a central AGN describes rather well the general properties of the optical emission line spectra ({\\em e.g.} Robinson et al. 1987). However, when the first high redshift radio galaxies (HZRG, $z>1.7$) were discovered the situation was not so clear. These objects showed very blue colours and strong Ly$\\alpha$ emission --- properties expected for galaxies in the process of formation ({\\em e.g.} Spinrad {\\em et al.} 1985, McCarthy {\\em et al.} 1987). It was proposed that some of these HZRG, like 3C326.1, were very young galaxies in which the strong Ly$\\alpha$ emission was powered by the young blue stars (McCarthy {\\em et al.} 1987). However, detailed analysis of the spectra showed that young stars could not explain the general properties of the UV emission lines of HZRG; in particular, the large emission line equivalent widths and the existence of highly ionized species, like CIV and HeII ({\\em e.g.} McCarthy {\\em et al.} 1990). In the context of the unified schemes for powerful radio galaxies ({\\em e.g.} Barthel 1989), we would expect these galaxies to harbour powerful AGN. This, together with the fact that AGN photoionization succeeded in explai\\-ning the optical emission line ratios of low redshift radio gala\\-xies, suggested that the same mechanism dominates the ionization processes in HZRG (McCarthy 1993). In consequence, most attempts to explain the UV emission line spectra of HZRG have invoked pure AGN photoionization. However, both imaging and spectroscopy of HZRG show that strong interactions are taking place between the advancing radio jet and the ISM of the host galaxy (section 2). Such interactions will generate powerful shocks which will disturb the morphology, kinematics and physical conditions (density, temperature, pressure) of the gas and potentially modify its ionization state. Currently a key issue in the study of these objects is the relative importance of jet-induced shocks and AGN illumination: do shocks dominate the emission line processes? Or is AGN photoionization dominant with the influence of the shocks mainly manifested in kinematic and morphological disturbances? We have addressed this problem by studying the UV emission line spectrum of a sample of very HZRG ($z>$1.7), comparing their line ratios with both shock and AGN photoionization models. We base our study on the UV lines because most of the information we have about HZRG is derived from studies of the UV emission line spectra. At such high redshifts the main optical diagnostic emission lines are shifted into the infrared, and most of the exis\\-ting IR spectra for the HZRG are of poor quality. In contrast, at low redshifts we have the opposite pro\\-blem: the optical diagnostic line ratios are well measured, but the UV spectra are of low quality. For the low redshift objects the optical line ratios have proved inefficient at distinguishing the ionization mechanism unambiguously, although it has been suggested that the UV line ratios might provide a stronger discriminant (Sutherland {\\em et al.} 1993). Therefore, the interest in understanding the emission of the UV lines can be extended to low redshift objects. We propose to develop a diagnostic method to discriminate between shock and AGN ionization in radio galaxies at all redshifts. We review in section 2 the observational evidence for AGN illumination and shocks in powerful radio galaxies. The data sample is described in section 3 and the diagnostic diagrams in section 4. Sections 5 and 6 present the models and the comparison with the data: section 5 concentrates on the effects of a) varying the shape of the AGN continuum and b) changing the excitation and ionization mechanism (shocks); section 6 analyzes the effects that different physical conditions in the extended gas can produce in the observed UV spectra. In section 7 we extend our diagnostic method to low redshift radio galaxies. Section 8 includes summary and conclusions. ", + "conclusions": "We have studied the UV line ratios of a sample of {\\it high} redshift radio galaxies to understand which is the main mechanism responsible for the line emission processes. The models which reproduce the optical line ratios of low redshift EELR (-1.5 power law or hot black body and solar abundances of the gas) cannot explain the UV line ratios of very high redshift radio galaxies. We have investigated several possibilities to explain these discrepancies. The most suggestive one, i.e. shocks produced in jet/cloud interactions, fails to explain the data. On the contrary, a harder ionizing AGN continuum (power law of index $\\alpha$=-1) and the classical -1.5 power law with a low metallicity produce an excellent fit to the data. The first possibility is supported by the existing evidence of an evolution towards harder continuum at higher redshifts of the mean spectral shape of the AGN conti\\-nuum. The second possibility is supported by the evidence provided by studies of absorption line systems in the line of sight of quasars, which demonstrate abundance of $\\sim$0.1 solar at very high redshifts. The UV diagnostic diagrams do not distinguish unambiguously between jet-induced shocks and AGN photoioni\\-zation for the HZRG with $high$ ionization state because the two sets of models overlap in the diagrams. We note, however, that not only are the sequences on the diagnostic diagrams better explained in terms of photoionization, but the overall ionization state of the HZRG is conside\\-rably higher than is measured in well-studied jet---cloud interactions at lower redshifts. Our result do not contradict the existence of shocks in these objects. Rather it suggests that the continuum emitted by the central AGN is the main ionizing source, with the shocks having an important effect on the kinematics and morphology of the gas. For the lower ionization states, the distinction between the shock and photoionization models is more clear-cut, and the UV diagnostic diagrams are a promising means of distinguishing the major ionization mechanism in the low redshift jet---cloud interaction candidates. Finally, we note that it is always dangerous to base strong conclusions about the ionization mechanisms on measurements of just a few emission lines in a particular spectral range. Our results are suggestive but not conclusive. It is now essential to check the consistency of the model results by obtaining spectra which cover diagnostics in {\\it both} the optical and UV for individual high and low redshift radio galaxies." + }, + "9701/astro-ph9701084_arXiv.txt": { + "abstract": "The correlation between the Mg absorption index and the velocity dispersion ($\\sigma$) of local elliptical galaxies is very tight. Because the Mg absorption depends on both metallicity and age of the underlying stellar population the observed Mg--$\\sigma$ relation constrains the possible variation in metallicity and age for a given velocity dispersion. For a time interval with no change in metallicity any variation of the Mg index is caused only by the aging of the stars. We have measured the Mg absorption and velocity dispersion of ellipticals in three clusters at a redshift of $z=0.37$ and established their Mg--$\\sigma$ relation. For any given $\\sigma$, the measured Mg absorption is weaker than the mean value for local ellipticals. Since the evolution of bright cluster ellipticals between $z=0.4$ and today is most probably only `passive' this reduction in Mg can be attributed solely to the younger age of the stellar population. The small weakening of the Mg absorption of the distant galaxies compared to the local values implies that most of the stars in cluster ellipticals must have formed at high redshifts ($z_f>2\\ldots 4$). The Mg--$\\sigma$ test is a very robust method to investigate the evolution of elliptical galaxies and has several advantages over traditional methods using luminosities. A remaining problem is the aperture correction necessary to calibrate observations of galaxies at different distances. Here, we show that our general conclusions about the epoch of formation still hold when aperture corrections are calculated assuming a dependence of the radial gradient of $\\sigma$ on the galaxy's effective radius rather than assuming no dependence as was done in all previous studies. ", + "introduction": "It is well known that all dynamically hot galaxies in the local universe follow the same linear relationship between the Mg absorption around $\\lambda_0\\approx5170$~\\AA\\ and the internal velocity dispersion $\\sigma$ (\\cite{DLBDFTW87,BBF93}). Although the galaxies span a wide range in Mg and $\\sigma$, the Mg--$\\sigma$ relation is very tight. A sample of luminous Coma ellipticals ($\\lg\\sigma\\ge2.3$), e. g., with Mg$_2\\in[0.25,0.36]$~mag has a standard deviation from the linear fit of only $\\sigma_{\\rm int}=0.011$~mag. Mg$_2$ as defined in the Lick system (\\cite{FFBG85}) comprises mainly the molecular absorption of MgH. Because the measurement of this index in redshifted galaxies is very noisy we use the atomic Mg$_b$ index instead. A linear transformation from Mg$_2$ to Mg$_b$ enables us to still use the 7~Samurai sample of Coma and Virgo ellipticals as the comparison at zero redshift. Both from observational (\\cite{Gonza93}) and theoretical data (stellar population synthesis of \\cite{Worth94}) we derived consistently: Mg$_b/{\\rm \\AA}=14.9\\pm0.5\\cdot{\\rm Mg}_2/{\\rm mag}$. In figure~\\ref{mgbs} small circles present the Mg$_b$--$\\sigma$ relation of the local comparison sample. A principal component analysis yields as best fit: \\begin{equation} \\label{gl_mgbs} \\mbox{Mg}_b = 2.7 \\lg \\sigma_0 - 1.65 \\end{equation} The dependence of Mg$_b$ on age and metallicity can be explored using stellar population synthesis models. From Worthey's 1994 models we derived the following formula that holds for ages $t>3$~Gyrs and metallicities $-2<\\lg Z/Z_\\odot<+0.25$: \\begin{equation} \\label{gl_mgbtZ} \\lg \\mbox{Mg}_b = 0.20 \\lg t + 0.31 \\lg Z/Z_\\odot + 0.37 \\end{equation} This formula allows to determine the maximum variation of both relative age and relative metallicity as it is constrained by the very tightness of the Mg$_b$--$\\sigma$ relation. For a given velocity dispersion $\\sigma$ (with $\\lg\\sigma\\ge2.3$) and zero variation in metallicity or age, resp., we find: \\begin{equation} \\label{gl_tZdisp} \\Delta t/t < 0.17 \\quad \\mbox{and} \\quad \\Delta Z/Z < 0.11 \\end{equation} This narrow constraint on the age spread of cluster ellipticals implies that they did not form continuously at the same rate but that there was a rather short formation epoch of these galaxies. If, e.g., the majority of ellipticals were formed 12~Gyrs ago, then the scatter in age would be about 2~Gyrs. However the formation epoch itself can only be determined by comparing the relative ages of galaxies at zero and a significantly higher redshift. ", + "conclusions": "first, the stellar population of elliptical cluster galaxies evolves predominantly passive between this redshift and today and second, the epoch of formation of the stars that mainly make up the ellipticals today is at high redshifts. The evolution in Mg$_b$ can be reliably transformed into an increase of the $B$--magnitude, thus, allowing the calibration of elliptical cluster galaxies as standard candles. \\bigskip {\\bf Acknowledgements:} This work is a collaboration with Drs. Bender, Belloni, Bruzual and Saglia and was supported by the ``Sonderforschungsbereich 375--95 f\\\"ur Astro--Teilchenphysik der Deutschen Forschungsgemeinschaft'' and by DARA grant 50 OR 9608 5." + }, + "9701/astro-ph9701037_arXiv.txt": { + "abstract": "We calculate the formation rates of low-mass X-ray binaries with a black hole. Both a semi-analytic and a more detailed model predict formation rates two orders of magnitude lower than derived from the observations. Solution of this conundrum requires either that stars with masses less than 20~\\msun\\ can evolve into a black hole, or that stellar wind from a member of a binary is accompanied by a much larger loss of angular momentum than hitherto assumed. ", + "introduction": "Six low-mass X-ray binaries currently are known to have a mass function that indicates a mass of the compact object larger than $\\sim 3 \\msun$, and these are therefore believed to have a black hole accretor. All these systems are soft X-ray transients and the total number in the galaxy of such systems is estimated to be between a few hundred and a few thousand; at any moment in time only a small fraction is X-ray active (for reviews, see \\cite{cow92}, \\cite{tl95}, \\cite{ts96}). The most popular formation scenario for low-mass X-ray binaries proposes a relatively wide binary with an extreme mass ratio as the progenitor system (\\cite{vdh83}). The massive star evolves to fill its Roche lobe and engulfs its low-mass companion. Due to the extreme mass ratio the low-mass star spirals into the envelope of the high-mass star. A close binary remains if the spiral-in ceases before the low-mass companion coalesces with the compact helium core of the primary. The helium core continues its evolution and may turn into a neutron star or a black hole. Only a few studies apply this evolutionary scenario specifically to the formation of low-mass X-ray binaries with an accreting black hole (see \\cite{rom92}). In this paper we discuss some problems of this standard formation scenario in its simplest form (Sect.~2) and show that these persist in a more refined treatment (Sect.~3). We quantify this in Sect.~4 and in Sect.~5 the results are discussed. ", + "conclusions": "Our computations, in contradiction to the results of \\cite*{rom92}, predict a formation rate for low mass X-ray binaries with a black hole which is much smaller than the value derived from the observed numbers and estimated X-ray lifetime. The striking difference between our and his results can be seen immediately by comparing his Fig.~1 with our Fig.~\\ref{Fig_rom92}. Romani's values for $a_{max}$ are much higher than ours: this is due to the fact that he uses $1/q$ instead of $q$ in the equation for the Roche lobe (our Eq.~\\ref{Roche_lobe}). (This error has been silently corrected in \\cite{rom94}.) This is true both for binaries in which mass transfer preceeds the supernova explosion, and for binaries in which the primary explodes after losing its envelope without ever having filled its Roche-lobe. The discrepancy between theoretical and observed formation rates cannot be solved by invoking different metallicities for the progenitor systems, nor by assuming different efficiencies for the envelope ejection during spiral-in. Black-hole binaries can be produced in larger numbers only if it is assumed that stars with initial masses less than approximately 20~\\msun\\ can collapse to black holes; or alternatively if it is assumed that the angular momentum loss caused by the stellar wind is so high that the binary orbit shrinks; or alternatively if the collapse of a helium core to a black hole is asymmetric, so that the post-supernova orbit can be smaller than the pre-supernova orbit (see \\cite{pzv96})." + }, + "9701/astro-ph9701201_arXiv.txt": { + "abstract": "We investigate the gas mass function of clusters of galaxies to measure the density fluctuation spectrum on cluster scales. The baryon abundance confined in rich clusters is computed from the gas mass function and compared with the mean baryon density in the universe which is predicted by the Big Bang Nucleosynthesis. This baryon fraction and the slope of the gas mass function put constraints on $\\sigma_8$, the rms linear fluctuation on scales of $8h^{-1}\\Mpc$, and the slope of the fluctuation spectrum, where $h$ is the Hubble constant in units of 100 $\\kms \\oMpc$. We find $\\sigma_8 = 0.80 \\pm 0.15$ and $n \\sim -1.5$ for $0.5 \\le h \\le 0.8$, where we assume that the density spectrum is approximated by a power law on cluster scales: $\\sigma(r) \\propto r^{-{3+n\\over{2}}}$. Our value of $\\sigma_8$ is independent of the density parameter, $\\Omega_0$, and thus we can estimate $\\Omega_0$ by combining $\\sigma_8$ obtained in this study with those from $\\Omega_0$-dependent analyses to date. We find that $\\sigma_8(\\Omega_0)$ derived from the cluster abundance such as the temperature function gives $\\Omega_0 \\sim 0.5$ while $\\sigma_8(\\Omega_0)$ measured from the peculiar velocity field of galaxies gives $\\Omega_0 \\sim 0.2-1$, depending on the technique used to analyze peculiar velocity data. Constraints are also derived for open, spatially flat, and tilted Cold Dark Matter models and for Cold $+$ Hot Dark Matter models. ", + "introduction": "To measure the spectrum of density fluctuations is one of the key issues in discussing the structure formation in the universe. Clusters of galaxies are suitable objects to measure the spectrum on scales of $\\sim 10 h^{-1}$ Mpc, where $h$ is the Hubble constant in units of 100 $\\kms \\oMpc$. This is because the abundance of clusters is sensitive to the nature of the spectrum, in particular the amplitude, and because the fluctuations on cluster scales can be reliably discussed by linear theory. Henry \\& Arnaud (1991) derived $\\sigma_8 = 0.59 \\pm 0.02$ and $n= -1.7^{-0.65}_{+0.35}$ from the X-ray temperature function of clusters for $\\Omega_0=1$ universes, where $\\sigma_8$ is the rms linear fluctuation on scales of $8h^{-1}\\Mpc$, $\\Omega_0$ is the cosmological density parameter, and they assumed that the density spectrum is approximated by a power law in wavenumber as $P(k) \\propto k^n$ on cluster scales. White, Efstathiou, \\& Frenk (1993a) obtained $\\sigma_8 \\simeq (0.57 \\pm 0.05) \\Omega_0^{-0.56}$ using the spatial number density of rich clusters. Similar results for $\\sigma_8$ were also obtained by other authors for open and spatially flat Cold Dark Matter (CDM) models using Henry \\& Arnaud's (1991) temperature function data; Eke, Cole, \\& Frenk (1996) found $\\sigma_8 = (0.50 \\pm 0.04) \\Omega_0^{-\\alpha}$ and Viana \\& Liddle (1996) found $\\sigma_8 = 0.6 \\Omega_0^{-\\alpha}$; in both estimates $\\alpha$ varies from $\\sim 0.4$ to $\\sim 0.6$ depending on the dark matter model assumed. There is, however, a problem with those measurements of $\\sigma_8$ that one cannot know $\\sigma_8$ unless $\\Omega_0$ is given; in general, the measured quantity is not $\\sigma_8$ but a combination of $\\sigma_8$ and $\\Omega_0$ like $\\sigma_8 \\Omega_0^{0.6}$. This problem is common to almost all methods for measuring $\\sigma_8$, such as the one using the peculiar velocity field of galaxies. In this paper, we measure observationally the fraction of baryons confined in clusters of galaxies to the total baryons in the universe using the cluster gas mass function, and give $\\Omega_0$-independent measurements of $\\sigma_8$ by comparing the observed baryon fraction with the theoretical prediction. As will be seen in \\S$\\ $3, the theoretical derivation of the baryon fraction does not depend on $\\Omega_0$, since the baryon fraction measures essentially the fraction of density fluctuations with an overdensity larger than the critical value. We find $\\sigma_8 = 0.80 \\pm 0.15$, with the quoted errors including the uncertainties in $h$ ($0.5 \\le h \\le 0.8$). The power-law index of the spectrum, $n$, is also measured. The plan of this paper is as follows. In \\S$\\ $2 we derive the cluster gas mass function from the observed sample and compute the fraction of baryons confined in rich clusters. In \\S$\\ $3, we describe an analytic model of gravitational halo formation to predict the fraction of mass confined in dark haloes with a given mass range. In \\S$\\ $4, we estimate $\\sigma_8$ and $n$. We also derive constraints on three sets of CDM models and a set of Cold $+$ Hot Dark Matter (CHDM) models in \\S$\\ $4. Finally, \\S$\\ $5 summarizes our conclusions. ", + "conclusions": "We have constructed observationally the gas mass function of clusters of galaxies to measure the fraction of baryons confined in clusters to the total baryons in the universe. Comparing this baryon fraction and the slope of the gas mass function with the prediction by the gravitational halo formation model, we have found $\\sigma_8 = 0.80 \\pm 0.15$ and $n \\sim -1.5$ for $0.5 \\le h \\le 0.8$. Our value of $\\sigma_8$ is independent of $\\Omega_0$, and thus we can estimate $\\Omega_0$ from the present result and previous ones in which $\\sigma_8$ was obtained as a function of $\\Omega_0$. We have found that $\\sigma_8(\\Omega_0)$ derived from the cluster abundance gives $\\Omega_0 \\sim 0.5$ while $\\sigma_8(\\Omega_0)$ measured from the peculiar velocity field of galaxies gives $\\Omega_0 \\sim 0.2-1$, depending on the technique used to analyze the peculiar velocity data. We have also examined four sets of cosmic structure formation models and have found that the following models match both the observed baryon fraction and the \\cobe data; open CDM models with $\\Omega_0 \\simeq 0.55-0.65$ (for $h=0.5$) and $\\Omega_0 \\simeq 0.35-0.45$ ($h=0.8$); spatially flat CDM models with $\\Omega_0 \\simeq 0.35 - 0.45$ ($h=0.5$) and $\\Omega_0 \\simeq 0.15 - 0.2$ ($h=0.8$); tilted CDM models with $n \\simeq 0.75-0.85$; CHDM models with $\\Omeganu \\gsim 0.15$." + }, + "9701/astro-ph9701171_arXiv.txt": { + "abstract": "Time profiles of gamma ray bursts (GRBs) are extremely diverse in their durations, morphologies, and complexity. Nevertheless, the average peak-aligned profile of all bursts detected by BATSE with sufficient data quality has a simple ``stretched'' exponential shape, $F \\propto \\exp[-(t/t_0)^{1/3}]$, where $t$ is the time measured from the time for the peak flux, $F_p$, of the event, and $t_0$ is a time constant. We study the behaviour of $t_0$ of both the post-peak and the pre-peak slopes of the average time profile as a function of the peak brightness range of the burst sample. We found that the post-peak slope shows time dilation when comparing bright and dim bursts, while the pre-peak slope hardly changes. Thus dimmer bursts have a different shape -- they are more asymmetric. This shape-brightness correlation is observed at a 99.6\\% confidence level. Such a correlation has a natural explanation within the pulse avalanche model, which is briefly described. Complex events, consisting of many pulses are more symmetric and are intrinsically brighter. Bursts consisting of one or a few pulses are intrinsically weaker and more asymmetric. For such a correlation to be observable requires that the luminosity distance distribution of GRBs to be different from a power-law. Keywords: Gamma-ray bursts, Methods: Data analysis. ", + "introduction": "Stern (1996) found that the average peak-aligned time profile of GRBs (the procedure of peak-alignment was pioneered by Mitrofanov et al. 1994, 1995) in the BATSE-2 catalog has a simple ``stretched'' exponential shape, \\\\ $=\\exp[-(t/t_0)^{1/3}]$, where $t$ is the time since the peak flux, $F_p$, of the event, and $t_0$ is a constant ranging from 0.3 s for strong bursts to $\\sim$ 1 s for dim bursts. This dependence of $t_0$ on brightness can be interpreted as a cosmological time dilation (e.g., \\cite{pac92,pir92}). Such a simple average time profile is remarkable considering the diverse and chaotic behavior of the individual time profiles of GRBs. On the other hand, the simple shape of the average time profile gives an excellent opportunity to study effects such as time dilation. Here, we study the two slopes of the average time profile for a larger sample of GRBs and with a more accurate treatment of the background than was done in Stern (1996). Another advantage we now have is the access to the pulse avalanche model developed by Stern \\& Svensson (1996) which successfully describes many statistical properties of GRBs including the stretched exponential shape of the profile and, of particular importance for the present work, the rms variance of individual time profiles. This means that we can rely on this model when estimating the errors of stretched exponential fits, which in turn gives us reliable estimates of the significance levels of the observed effects. ", + "conclusions": "Besides the time dilation effect, observed in many previous works (e.g., Norris et al. 1994) we also see a dependence of the profile asymmetry on brightness and this effect is of the same order of magnitude as the time dilation itself. A correlation of such a kind cannot be due to spectral redshifts. We find that strong events have slightly smaller asymmetry in the higher energy bands, (LAD channels 4 and 3) than in the lower energy bands (LAD channels 1 and 2). Redshifting the softer part of the spectrum below the detector threshold would then give rise to more symmetric, rather than asymmetric profiles. Details will be published elsewhere. The effect of trigger efficiency is negligible for the brightness groups considered here. Also one can hardly find an evolutionary factor that would change the asymmetry. One can suggest that there exist two separate classes of GRBs with different degree of asymmetry, which are differently distributed in space (separate classes of {\\it long} bursts are required as short bursts do not contribute to the asymmetry). But such a suggestion seems too arbitrary, too radical, and unnecessary as there exists a much simpler explanation. The simplest explanation is that the observed correlation is a consequence of an intrinsic correlation between shape and brightness as described above. The necessary condition for such a correlation to be observable is a significant deviation from a power law for the GRB distribution over luminosity distance. This would allow intrinsically strong events to dominate in the brightest observational range. The observed log $N$ - log $P$ distribution is actually curved (Meegan et al. 1996) and this is natural if the distance distribution covers both Euclidean and $z \\sim 1$ regions which have different luminosity distance scalings. Maybe our detection of a strong shape - brightness correlation imposes a stronger constraint on the curvature of the true radial distribution of GRBs than what follows from the observed log $N$ - log $P$ distribution. Detailed studies are, however, required to formulate this intuitive conclusion at a quantitative level. As one kind of correlation has been observed, other kinds of intrinsic correlations may also be observable and this causes a problem for the cosmological interpretation of the time dilation effect. This problem, considered by Brainerd (1994), arises from unavoidable correlations between peak luminosity and time scales caused by different bulk Lorentz factor in the sources of different bursts. This effect can mimic both cosmological time dilation and spectral redshift. However, if our interpretation using the pulse avalanche model is valid we must conclude that the real time dilation is larger than that obtained from Table 1. Actually, within the pulse avalanche framework, intrinsically weak events are not only more asymmetric, but they are also narrower (see column $t_r+t_d$ in Table 2). This is a new correction that increases the real time dilation more than the correction arising from spectral redshift (see Norris et al. 1994). The corrected time dilation could exceed a factor 2 and it could be caused by different effects, including the cosmological one. Unfortunately, the task of extracting the cosmological component from the total time dilation seems extremely difficult." + }, + "9701/astro-ph9701059_arXiv.txt": { + "abstract": "First results of the inversion of Stokes I and V profiles from plage regions near disk center are presented. Both low and high spatial resolution spectra of \\ion{Fe}{1} 6301.5 and \\ion{Fe}{1} 6302.5 \\AA\\ obtained with the Advanced Stokes Polarimeter (ASP) have been considered for analysis. The thin flux tube approximation, implemented in an LTE inversion code based on response functions, is used to describe unresolved magnetic elements. The code allows the simultaneous and consistent inference of all atmospheric quantities determining the radiative transfer with the sole assumption of hydrostatic equilibrium. By considering velocity gradients within the tubes we are able to match the full ASP Stokes profiles. The magnetic atmospheres derived from the inversion are characterized by the absence of significant motions in high layers and strong velocity gradients in deeper layers. These are essential to reproduce the asymmetries of the observed profiles. Our scenario predicts a shift of the Stokes V zero-crossing wavelengths which is indeed present in observations made with the Fourier Transform Spectrometer. ", + "introduction": "A remarkable, still unexplained feature of the spectra emerging from facular and network regions in the solar photosphere is the conspicuous asymmetry exhibited by Stokes V profiles. At disk center, the peak and absolute area of the blue lobe are larger than those of the red lobe for the majority of \\ion{Fe}{1} lines (Solanki \\& Stenflo 1984). In addition, the red lobe has a more extended wing than its blue counterpart. A detailed analysis of the properties of high spatial resolution Stokes V spectra in plage regions can be found in Mart\\'{\\i}nez Pillet et al.\\ (1997). The search of the processes that give rise to the area asymmetry of Stokes V has been one of the most vivid discussions of recent times in solar physics. Illing et al.\\ (1975) first suggested that gradients of velocity and magnetic field along the line of sight may produce asymmetric V profiles. Originally, this mechanism was put forward to explain broad-band measurements of circular polarization in sunspots. Later, a series of papers refined this idea and settled down the physics involved (van Ballegooijen 1985; S\\'anchez Almeida et al.\\ 1988; Grossmann-Doerth et al.\\ 1988). In the current picture, the area asymmetry is produced by the combined (but otherwise {\\em spatially separated}) gradients of magnetic field and velocity that photons traversing the tubes encounter at the boundary layer. These gradients are generated by the expanding walls of magnetic elements embedded in field-free surroundings. The canopy mechanism, however, does fail to explain the peak asymmetry of Stokes V spectra. Observationally, the peak asymmetry turns out to be several times larger than the area asymmetry. All attempts to reproduce the observations have invariably led to ratios of peak to area asymmetries of the order of unity. A number of other scenarios have been invoked to solve this problem. They include NLTE effects (Kemp et al.\\ 1984; Landi Degl'Innocenti 1985), linear oscillations (Solanki 1989) or longitudinal waves (Solanki \\& Roberts 1992) within the magnetic elements, non-linear oscillations (Grossmann-Doerth et al.\\ 1991) and, more recently, micro-structured magnetic atmospheres (S\\'anchez Almeida et al.\\ 1996). None of these mechanisms has proven yet to be able to reproduce the observations. Non-linear, high-frequency oscillations inside the tubes have never been detected in spite of the observational efforts. Time series of Stokes V profiles at disk center analyzed by Fleck (1991), for example, did not reveal such motions. Micro-structured magnetic atmospheres, on the other hand, seem promising but still need further development. The present state-of-the-art can be summarized as follows: no physical process is known to be able to generate a peak asymmetry significantly larger than the area asymmetry exhibited by Stokes V spectra in plage and network regions at disk center. In fact, some of the mechanisms above were proposed, without much success, as a very last attempt to explain the observations. It is our belief, however, that velocity gradients have not been completely exploited yet. A natural extension of the canopy mechanism might include mass motions in the magnetized plasma. Empirical evidence of such motions has been found repeatedly (Solanki 1986, 1989; S\\'anchez Almeida et al.\\ 1990), but stationary flows were ruled out already in the earliest analyses as the dominant source of the asymmetries because they would generate Doppler shifts of the Stokes V zero-crossing wavelengths larger than the reported upper limit of $\\pm 250$ m s$^{-1}$ (Solanki 1986). Some preliminary calculations, often involving one dimensional models, have confirmed this result (e.g., Solanki \\& Pahlke 1988). However, it could be possible that the small shifts of the zero-crossings are effectively produced by velocity gradients in the magnetic atmosphere. These gradients might explain the marginal dependence between the observed shifts and line strength that various analyses have revealed (Solanki 1986; S\\'anchez Almeida et al.\\ 1989). If this is the case, stationary flows should be reconsidered again as a possible source of the asymmetries. In order to proceed further, reliable 2-D fluxtube atmospheres have to be used. Unfortunately, no such models exist. Numerical simulations (e.g.\\ Steiner et al.\\ 1995) are more intended to understand the interaction between magnetic elements and their convective surroundings than to explain the observed spectra. For this reason it is not strange that, despite their important results, they are not able to reproduce the asymmetries. On the other hand, empirical models fail to match the observations because they neglect the role of velocity gradients to diminish the complexity of the problem. Indeed, extracting the information contained in the Stokes spectrum is hampered by the intricate non-linear dependences of the radiative transfer equation on the various quantities defining both the thermodynamical and magnetic properties of the atmospheres. The trial-and-error method does not work in this case, and so inversion techniques come into play. Here we follow an empirical approach to investigate the capabilities of velocity-based mechanisms. As a first step, a new LTE inversion code of the radiative transfer equation particularized to the case of thin flux tubes has been developed. Applied to real observations, it carries out a simultaneous inference of the whole set of model parameters which reproduce the observed Stokes spectrum, thus ensuring self-consistency. The details of the procedure and numerical tests will be presented in a forthcoming paper, although a brief description of the method is given below. This letter reports on fluxtube model atmospheres derived from the inversion of high spatial ($\\sim 1\"$) and temporal ($\\sim 4$ sec) resolution Stokes I and V spectra of \\ion{Fe}{1} 6301.5 and 6302.5 \\AA\\ from plage regions at disk center obtained with the Advanced Stokes Polarimeter (ASP, Mart\\'{\\i}nez Pillet et al.\\ 1997). For the first time, plausible model atmospheres are found that reproduce the whole shape of the ASP \\ion{Fe}{1} 6301.5 and \\ion{Fe}{1} 6302.5 \\AA\\ Stokes spectra to a degree of accuracy never reached before. The recovered model atmospheres, in which stationary flows within the tubes were allowed, have been used to synthesize the Stokes profiles of a large number of spectral lines. The comparison of their zero-crossing wavelengths to observations made with the Fourier Transform Spectrometer (FTS, Stenflo et al.\\ 1984) reveals that the observed Doppler shifts are not randomly distributed and suggests that non-negligible mass motions within the magnetic elements do exist. ", + "conclusions": "" + }, + "9701/astro-ph9701090_arXiv.txt": { + "abstract": " ", + "introduction": "One of the goals of cosmology is to understand the large scale structure of the universe. Gravitational instability models hold that large-scale structure forms as the result of gravitational amplification of initially small perturbations in the primordial density distribution. The initial ``seeds'' may result from quantum fluctuations in a scalar field (inflation) or causal ordering following a phase transition with broken symmetry (topological defects), but in all models the evolution occurs in a ``background'' cosmology described by the Friedmann-Robertson-Walker (FRW) metric. The anisotropy of the cosmic microwave background (CMB) provides an observational test of the the assumption that on large scales spacetime asymptotically approaches the FRW metric. The detection of fluctuations in the CMB provides support for a nearly homogeneous spacetime, limiting density perturbations in the early universe to the level $\\delta \\rho / \\rho \\approx 10^{-5}$. The CMB anisotropy also provides a direct test of the assumption of isotropy about each spatial point. Small deviations from the FRW metric lead to observable signatures in the CMB. Open or flat models with global rotation or shear will exhibit a spiral pattern of temperature anisotropy resulting from the handedness of the geodesics propagating through an anisotropic spacetime. Further geodesic focusing creates ``hot spots'' in open models (density $\\Omega_0 < 1$), while closed models exhibit a pure quadrupole pattern. Several authors have used the CMB to limit rotation $\\omega$ and shear $\\sigma$ in the universe \\cite{collins_1973,barrow_1985,smoot_1993,bunn_1996}. Prior to the detection of CMB anisotropy, Collins \\& Hawking and Barrow, Juskiewicz, \\& Sonoda used upper limits on the CMB quadrupole amplitude to limit $\\omega/H_0 < 10^{-5}$ for flat or moderately open ($\\Omega_0 \\approx 0.3$) models. Smoot used the quadrupole detection from the first-year {\\it COBE} DMR sky maps to limit $\\omega/H_0 < 10^{-6}$. Recently, Bunn, Ferreira, \\& Silk fitted the full spiral pattern from models with global rotation to the 4-year DMR data and derived limits $\\omega/H_0 < 3 \\times 10^{-7}$. The DMR anisotropy data are dominated by a power-law spectrum of fluctuations which are ill-described by the pattern of angular anisotropy predicted for models with global rotation or shear. Limits to rotation or shear based on full-sky ``template'' maps must account for the presence of this power-law component. Bunn, Ferreira, \\& Silk \\cite{bunn_1996} use a least-squares fit including only instrument noise, and then account for chance alignment of the spiral template map with features in the power-law component using Monte Carlo simulations. In this paper, we present an independent analysis using a formalism which expressly includes the dominant power-law component in the fitting process, and derive upper limits $\\omega/H_0 < 6 \\times 10^{-8}$ for $0.1 \\le \\Omega_0 \\le 1$. ", + "conclusions": "" + }, + "9701/astro-ph9701133_arXiv.txt": { + "abstract": "The temperature and abundance structure in the intracluster medium (ICM) of the Hydra-A cluster of galaxies is studied with {\\it ASCA} and {\\it ROSAT}. The effect of the large extended outskirts in the point-spread function of the X-Ray Telescope on {\\it ASCA} is included in this analysis. In the X-ray brightness profile, the strong central excess above a single $\\beta$-model, identified in the {\\it Einstein} and {\\it ROSAT} data, is also found in the harder energy band ($>$4~keV). A simultaneous fit of five annular spectra taken with the GIS instrument shows a radial distribution of the temperature and metal abundance. A significant central enhancement in the abundance distribution is found, while the temperature profile suggests that the ICM is approximately isothermal with the temperature of $\\sim$3.5~keV. The {\\it ROSAT} PSPC spectrum in the central $1'.5$ region indicates a significantly lower temperature than the GIS result. A joint analysis of the GIS and PSPC data reveals that the spectra can be described by a two temperature model as well as by a cooling flow model. In both cases, the hot phase gas with the temperature of $\\sim$3.5~keV occupies more than 90\\% of the total emission measure within $1'.5$ from the cluster center. The estimated mass of the cooler (0.5--0.7~keV) component is $\\sim$2--6$\\times 10^9$ M$_{\\odot}$, which is comparable to the mass of hot halos seen in non-cD ellipticals. The cooling flow model gives the mass deposition rate of $60 \\pm 30$ M$_{\\odot}$ yr$^{-1}$, an order of magnitude lower than the previous estimation. ", + "introduction": "X-ray imaging studies with the {\\it Einstein} observatory and the {\\it ROSAT} satellite revealed that many clusters exhibit central concentrations in their X-ray brightness profiles. The central excess emission is often interpreted as a result of a cooling flow, a thermal instability occurring at the densest part of cluster (see \\cite{Fabian94} for a review). Edge et al. (1992) reported that about two thirds of all known clusters have such structure, suggesting that this structure is a fairly common characteristic among clusters. Spectroscopic studies with {\\it Einstein} and {\\it ROSAT} found that there are cooler gas components with temperature $\\sim$10$^7$~K in cluster centers (e.g. \\cite{Cani79}). However, the previous spectroscopy is limited to the soft X-ray band ($\\sim$0.1--4~keV). {\\it ASCA} has provided the first opportunity to resolve X-ray spectra spatially in the wider 0.5--10~keV energy range. In this paper we report the {\\it ASCA} observation of the Hydra-A cluster of galaxies ($z=0.0522$). Its X-ray luminosity measured with the {\\it Einstein} Observatory was 4.1$\\times$10$^{44}$ ergs~s$^{-1}$ (0.5--4.5~keV) (\\cite{David90}), and a consistent value of 4.8$\\times$10$^{44}$ ergs~s$^{-1}$ in the same energy band was indicated from the {\\it Ginga} observation which measured the X-ray spectrum in the 2--10~keV band (\\cite{Tsuru}). These luminosities are among the largest of poor clusters (\\cite{Kriss83}; \\cite{Tsuru}). The average X-ray temperature of $\\sim$4~keV, as measured with {\\it Ginga}, is also rather high for a poor cluster. The central region of this cluster is of particular interest. The cD galaxy 3C218 = Hydra-A is a strong radio source, with a very complex radio morphology and the highest Faraday rotation ever measured from a radio galaxy (\\cite{Ekers}; \\cite{Kato}; \\cite{Taylor}). Furthermore, the X-ray surface brightness profile obtained with the {\\it Einstein} IPC exhibits a large central excess deviating from an isothermal $\\beta$-model, suggesting the presence of a cooling flow with an estimated mass deposition rate of 600 $\\pm$ 120 $M_{\\odot}$ yr$^{-1}$ (\\cite{David90}). This is one of the largest values thus far attributed to the X-ray emitting gas in clusters of galaxies (\\cite{Edge92}). When we analyze the {\\it ASCA} data, the complex response of the X-Ray Telescope (XRT; \\cite{XRT}) onboard {\\it ASCA} introduces considerable difficulty for spatially resolved spectroscopy. To confront with these difficulties, some analysis methods have been proposed (\\cite{IkebeD}; \\cite{MEA96}; \\cite{Churazov96}). The method employed in this paper is based on the similar idea in Markevitch et al (1996) but is performed with different implementation. Since the position resolution of {\\it ASCA} is relatively poor, the sharp {\\it ROSAT} image is also very helpful to analyze the {\\it ASCA} data. Assisted by the {\\it ROSAT} image, we measured the spatial distribution of the temperature and metal abundance in the Hydra-A cluster from the {\\it ASCA} observation. In section 2, we describe the {\\it ASCA} observation and the data selection procedure. Section 3 gives the results using a conventional analysis method on the GIS and SIS data. Section 4 describes the analysis method developed here. In section 5, we analyze the X-ray brightness profiles in different energy bands. In section 6, we describe the temperature and abundance structure studied using the GIS spectra, and we also discuss the joint analysis of the {\\it ROSAT} and {\\it ASCA} data to study the temperature structure in the central region. We discuss the results and summarize them in section 7. $H_0 = 50\\ {\\rm km\\ s^{-1}\\ Mpc^{-1}}$ is assumed throughout. ", + "conclusions": "The energy-dependent flux-mixing effect due to the XRT PSF makes it quite difficult to analyze the data of extended sources observed with {\\it ASCA}. In order to perform simultaneous fitting of the X-ray images and spectra, we have calibrated the XRT PSF using the data from Cyg~X-1, and have fully taken it into account in the data analysis. The X-ray radial brightness profile obtained by the {\\it ROSAT} PSPC exhibits a central excess above a single $\\beta$-model within $\\sim$$1'.5$, as does the radial profile from the {\\it Einstein} IPC data. With {\\it ASCA}, we observed the X-ray surface brightness in higher energy bands up to $\\sim$10~keV for the first time. Using the newly developed analysis technique, we fit the radial brightness profiles and found that there is no clear difference among the profiles in different energy bands. This suggests that the central excess emission found in the soft X-ray images from {\\it Einstein} and {\\it ROSAT} also exists in the higher energy band of 4--10~keV. In particular, we successfully reproduced the GIS radial brightness profile as a sum of two $\\beta$-models. The narrower of the two $\\beta$ components, thought to represent the central excess emission, is required not only below 2~keV but also in harder energy bands, up to 10~keV, assuming a double $\\beta$ profile in all energy bands. The simultaneous fitting of the five annular spectra taken from the GIS data gave the radial profiles of the temperature and metal abundance (see Fig.~\\ref{gis-doujifit-fig}). The obtained overall temperature structure is consistent with being isothermal; this result is also consistent with the fact that all the radial profiles are very similar. However, the PSPC spectra accumulated from the region within $1'.5$, where the surface brightness profile begin to deviate from the $\\beta$-model, gives a significantly lower temperature than that obtained from the GIS spectrum. This means that there must be an additional cool component at the cluster center. Therefore we jointly analyzed the GIS and PSPC data and successfully fit both spectra simultaneously with the two temperature model as well as the cooling flow model (Fig~\\ref{2temp-model-fig}). The cooling flow model gives the mass deposition rate of $60\\pm30$ M$_{\\odot}$ yr$^{-1}$, an order of magnitude smaller than the 600 M$_{\\odot}$ yr$^{-1}$ estimated from the {\\it Einstein} data by David et al. (1990). If we assumed that all the flux coming from the central $1'.5$ region is originated from cooling flow, we would derive a value consistent with 600 M$_{\\odot}$ yr$^{-1}$, using the formula of $\\dot{M} = 2\\mu m_p L_x / 5kT$ and the bolometric luminosity within $1'.5$ of $L_x \\sim 4 \\times 10^{44}$ ergs s$^{-1}$. However, as we showed in \\S~6, the central $1'.5$ region can not be entirely cooled. More than 90\\% of the total emission measure consists of the hot non-cooled component with the temperature of $\\sim$3.5~keV. Therefore, the central excess in the X-ray brightness profile can not be formed only by the cool component. Since the central region representing the excess emission mostly occupied by the same hot ICM component which permeates the rest of the cluster, we interpret the central excess emission as an evidence for gravitational potential structure. The potential structure can be interpreted as consisting of two distinct components; a large scale cluster component and a central compact component attributed to the cD galaxy. In previous investigations, such a dual potential structure has been suggested (\\cite{TFN87}; \\cite{NB95}) or assumed (e.g. \\cite{Stewart}). The first direct observational evidence of the additional potential dimple around NGC~1399 in the Fornax cluster was found from the {\\it ASCA} observation (\\cite{Ikebe96}). In the case of the Hydra-A cluster, there also must be a central potential dimple around the cD galaxy which primarily causes the central excess in the brightness profile. Thus far, the central excess brightness seen in many clusters has been interpreted mainly as due to the central temperature decrease of gas, thus providing a basis for the cooling flow. However, our results clearly reveal that the central excess brightness is at least partially caused by the dual potential structure around the cD galaxy. Therefore, the cooling flow rate derived from the central excess brightness can be grossly over-estimated, as in the case of the Hydra-A cluster, if the dual potential structure is ignored. The metal abundance distribution is also a very important subject which is strongly related to the cluster evolution scenario. A detailed measurement has become possible for the first time using {\\it ASCA}. In the Hydra-A cluster, we found an indication of a central concentration in the metal abundance distribution (Fig.~\\ref{gis-doujifit-fig}). {\\it ASCA} has clearly detected the central concentration in the heavy element abundance distributions in the nearest two clusters, Virgo and Centaurus (\\cite{Matsumoto}; \\cite{Fukazawa94}), which are the firm confirmation of the previous results by {\\it Ginga} (\\cite{Koyama90}) and {\\it ROSAT} (\\cite{AF94}), respectively. A similar feature has also been discovered in the poor cluster AWM7 (\\cite{Xu}). The central abundance concentration may be caused by a large contribution by metal-enriched ISM of the cD galaxies, which may have not suffered a ram-pressure-stripping process in the cluster evolution (\\cite{Tamura96}), because the cD galaxy is sitting in the bottom of the gravitational potential well. We speculate that all the clusters showing central excess emission in the brightness profile have central concentration of heavy elements. Based on the results from the simultaneous fitting of the GIS and PSPC spectra with the two temperature model (model 1 in \\S~6.2), we calculated the mass of the cool component gas. We assumed that the cool component coexists with the hot component within the $1'.5$ region, and the local pressure balance is achieved between the cool component and the hot component as $n_{cool} T_{cool} = n_{hot} T_{hot}$. The mass of the cool component gas is estimated to be $M_{gas} \\sim 6 \\times 10^9$ M$_{\\odot}$; and its filling factor, that is, the volume fraction of the cool component, is $\\sim$6 $\\times 10^{-4}$. On the other hand, if the cool component is concentrated at the center, and the pressure equilibrium is achieved at the boundary of the cool component and the hot component, the cool component would be distributed out to $\\sim$5~kpc and would have a total mass of $M_{gas} \\sim 2 \\times 10^9$ M$_{\\odot}$. These values are about 0.1--0.4\\% of the stellar mass of the cD galaxy 3C218 = Hydra-A, estimated to be $1.5 \\times 10^{12}M_{\\odot}$ based on the assumption of $M/L_B=6(M/L_B)_{\\odot}$ and log$L_B = 11.4$. This is comparable to that in other non-cD ellipticals (\\cite{FJT})." + }, + "9701/astro-ph9701075_arXiv.txt": { + "abstract": "A new methanol maser line $6_{-1}-5_0~E$ at 133~GHz was detected with the 12--m Kitt Peak radio telescope using remote observation mode from Moscow. Moderately strong, narrow maser lines were found in DR21(OH), DR21--W, OMC--2, M8E, NGC2264, L379, W33--Met. The masers have similar spectral features in other transitions of methanol$-E$ at 36 and 84~GHz, and in transitions of methanol$-A$ at 44 and 95~GHz. All these are Class~I transitions, and the new masers also belong to Class~I. In two other methanol transitions near 133~GHz, $5_{-2}-6_{-1}~E$ and $6_2-7_1~A^+$, only thermal emission was detected in some sources. Several other sources with wider lines in the transition $6_{-1}-5_0~E$ also may be masers, since they do not show any emission at the two other methanol transitons near 133~GHz. These are NGC2071, S231, S255, GGD27, also known as Class~I masers. The ratio of intensities and line widths of the 133~GHz masers and 44~GHz masers is consistent with the saturated maser model, in which the line rebroadening with respect to unsaturated masers is suppressed by cross relaxation due to elastic collisions. ", + "introduction": "The widespread galactic methanol masers were found in many transitions in the microwave band. Two kinds of sites are associated with methanol masers. Cold dust cores sometimes containing star formation regions at a very early, pre--stellar stage of evolution are associated with Class I methanol masers. They emit in a particular set of methanol transitions at 25, 36, 44, 84, and 95~GHz. Another kind of star formation regions with young stars which are embedded in nascent molecular cores and produce ultra--compact HII regions, is associated with methanol masers of Class~II. They emit lines in a different from Class~I set of transitions, typically, at frequencies 6.7 and 12~GHz. Recently a new Class~II methanol maser line was detected at a very high frequency 157~GHz (Slysh et al. 1995) (see Menten (1991) for an earlier discussion of methanol maser classification). Masers of Class~I in methanol--$E$ were found in transitions between the levels of the $K=-1$ stack and those of the $K=0$ stack. Zuckerman et al. (1972) proposed a model of maser inversion by collisional excitation of molecules, followed by radiative decay. A much faster spontaneous decay of the $K=0$ levels leads to their underpopulation relative to the $K=-1$ levels, which results in the inversion of the transitions between the $K=-1$ and $K=0$ levels. Masers are expected in transitions $J_{-1}-(J-1)_0~E$ with $J\\ge 4$, when the $J_{-1}$ levels become higher than the $(J-1)_0$ levels. The first transition of this type is $4_{-1}-3_0~E$, at 36 GHz, and was discovered as a maser by Morimoto et al. (1985) in Sgr B2. The second transition, $5_{-1}-4_0~E$, was discovered at 84~GHz as a strong maser by Batrla and Menten (1988) in DR21(OH). The next transition is $6_{-1}-5_0~E$, at 133~GHz. The methanol line in this transition was first observed by Cummins et al. (1986) from Sgr B2, but it was not clear if it was a maser line. In this paper, we report on the detection of the maser emission from many sources in the transition $6_{-1}-5_0~E$ at 133~GHz. ", + "conclusions": "The similarity between line profiles of the new 133~GHz methanol masers and methanol masers of Class~I as well as association with the same sources, suggests that the 133~GHz methanol masers belong to the Class~I. The corresponding methanol transition $6_{-1}-5_0~E$ is of the same type $K=-1$ $\\to$ $K=0$, as most of the Class~I transitions in methanol--$E$. The detection of new masers in the transition $6_{-1}-5_0~E$ supports the suggestion that $K=-1$ ladder levels are overpopulated relative to $K=0$ ladder levels, which is the cause for the inversion of $J_{-1}-(J-1)_0$ transitions. A natural way of overpopulating the $K=-1$ ladder is collisional excitation followed by spontaneous decay. The decay rates are different for the two ladders: for the upper $6_{-1}$ level the spontaneous decay rate is 1.07$\\times$10$^{-4}$~s$^{-1}$, while for the lower $5_0$ level it is a factor of two higher 2.12$\\times$10$^{-4}$~s$^{-1}$. The molecules excited by collisions to the $6_{-1}$ level remain there longer than at the $5_0$ level, resulting in the higher population of the $6_{-1}$ level. The higher population of the upper level in the $6_{-1}-5_0~E$ transition, or population inversion (relative to the normally lower population of the upper level in thermally excited molecules) may produce the maser emission if the column density is high enough to provide significant amplification. The same mechanism of level population inversion holds for other similar transitions in methanol--$E$: $4_{-1}-3_0$, $5_{-1}-4_0$, $7_{-1}-6_0$, $8_{-1}-7_0$, and so on. In methanol--A, the $7_0-6_1$, $8_0-7_1$, $9_0-8_1$ etc. transitions can be inverted in the same manner. A comparison of 133~GHz masers with the strongest Class~I methanol masers at 44~GHz shows that both show emission features at the same radial velocities. One can suppose that the emission features at both frequencies originate from the same region. Compared to the 44~GHz masers, the 133~GHz masers are weaker: the flux density is a factor of 5 to 27 lower, and the line width is a factor of 1.4 to 4 larger. The flux density ratio is roughly consistent with the calculated optical depth ratio of about 6; this is expected for saturated masers at both frequencies. On the other hand, the observed larger line width at 133~GHz as compared to 44~GHz is more appropriate for unsaturated masers, where the line width is inversely proportional to the square root of the optical depth. This apparent contradiction can be resolved if the masers are saturated but the line rebroadening due to the saturation does not take place, because the velocity distribution undergoes cross relaxation toward a Maxwellian distribution as a result of elastic collisions (Nedoluha and Watson 1988). Therefore, line narrowing can be expected in the saturated masers, also; in the case of Class~I methanol masers, collisions are important both for the maser pump and for the velocity distribution relaxation, which permits line narrowing." + }, + "9701/astro-ph9701061_arXiv.txt": { + "abstract": "Line asymmetries and shifts are a powerful tool for studying velocity fields in the stellar photospheres. Other effects, however, could also generate asymmetries blurring the information of the velocity patterns. We have studied the shifts and asymmetries induced in the profiles of spectral lines by pressure effects. The best theoretical and experimental data on line broadening and shifts caused by collisions with atomic hydrogen were used to analyze the Na\\,{\\sc i} D and three Ca\\,{\\sc i} lines. Line bisectors of synthetic spectra computed with accurate data for the Na\\,{\\sc i} and Ca\\,{\\sc i} lines are compared with very high resolution high signal-to-noise ratio solar spectra and indicate that pressure broadening reproduces the wings of the observed lines, but pressure shifts introduce neither asymmetries nor shifts comparable to the observed ones. ", + "introduction": "\\label{sec1} In the study of stellar atmospheres, new insight is coming from the very high resolution spectroscopic techniques. Important physical information is contained in the spectral line asymmetries. They are supposed to be due to velocity fields in the line formation region and, in particular for late-type stars, they are mainly associated with solar-like granular motions (\\markcite{gray80}Gray 1980; \\markcite{dravins81}Dravins, Lindegren, \\& Norlund 1981). In metal-poor stars, with atmospheres in which the lower opacity could be responsible for an enhanced convective influence on the line formation, stronger deviations from symmetry are expected. \\markcite{viena}Allende Prieto et al. (1995) observed with very high spectral resolution ($\\lambda /\\Delta\\lambda\\sim 170000$) a set of isolated lines in the star HD140283, with a metallicity $\\sim 500$ times lower than that of the Sun. They found that the bisectors of the lines show the same kind of ``C'' shape observed in late-type metal-rich stars but with a more pronounced shift to the red in the wings. Those observations are part of a wider project aimed at better understanding the atmospheres of metal-poor stars. Velocity gradients and granulation patterns are, however, not the only potential factors that can produce spectral line asymmetries. Blends with other lines may occur, although they will be less important for metal-poor stars than for stars with higher metallicity. Isotopic shifts and hyperfine structure will split the line profile in different components, and introduce asymmetry (\\markcite{kuruczisot93}Kurucz 1993). On the other hand, it is well-known that collisions between an emitting/absorbing atom in a plasma and the surrounding particles induce not only a symmetric broadening (\\markcite{victor32}Weisskopf 1932), but also a shift in the line emission/absorption coefficient (\\markcite{lindholm42}Lindholm 1942,\\markcite{lindholm45}1945; \\markcite{foley46}Foley 1946). That displacement is proportional to the density of perturbers. The quantum-mechanical reformulations (\\markcite{anderson52}Anderson 1952, \\markcite{griem62}Griem et al. 1962) do not qualitatively change the classical result. When applied to stellar atmospheres, since the line formation takes place along a density stratified region, asymmetries and a net shift in the central wavelength of the emergent profile will appear. These line shifts were suggested to be the origin for the solar limb effect (\\markcite{spitzer50}Spitzer 1950; \\markcite{hart74}Hart 1974), although later the convective blue-shift, resulting from the correlation between intensity and velocity fields in the solar photosphere, became a better candidate (\\markcite{BeckVeg79}Beckers \\& de Vegvar 1978; \\markcite{BeckNelson78}Beckers \\& Nelson 1978). In these previous studies, making use of the Lennard-Jones potential (\\markcite{hindmarsh67}Hindmarsh, Petford, \\& Smith 1967) and the classical impact theory, pressure shifts in the line absorption profile were estimated at a representative atmospheric height only. This provides an indication of the mean displacement of the line but does not give information about possible asymmetries in its shape. The asymmetries will depend on the variation of the pressure shifts from the top to the bottom of the line formation region. \\markcite{vince87}Vince \\& Dimitrijevi\\'{c} (1987) analyzed the influence of pressure shifts in the spectral line formation making use of the Roueff theory (\\markcite{roueff70}Roueff 1970), and found that the effect might be responsible for up to $20-30$\\% of the asymmetry observed in the Na\\,{\\sc i} $\\lambda$ 6160.8 \\AA\\ solar line. For a better understanding of the behaviour of line asymmetries in the atmospheres of late-type stars, we analyze in this paper the potential asymmetries present in the line profile due to collisional shift gradients. The study is carried out for a few atomic lines which have been the object of the best theoretical and experimental studies on line broadening and shifts caused by collisions with atomic hydrogen. We have concentrated mainly on the Sun, for which observed spectra with the highest resolution and signal-to-noise ratio are available. In section 2 we estimate the asymmetries in the spectral profiles of Na\\,{\\sc i} and Ca\\,{\\sc i} lines emerging from the solar atmosphere. We also estimate the shifts for a given Ca\\,{\\sc i} line in a metal-poor star and compare them with our observations. The main conclusions are summarized in section 3. ", + "conclusions": "\\label{sec4} We have analyzed carefully the asymmetries associated with pressure effects in a limited sample of spectral lines. We used the most accurate experimental and theoretical data available for Na D and three Ca\\,{\\sc i} lines. All of the computed synthetic spectra were compared with their corresponding counterparts in the solar spectrum. The predictions for the Ca\\,{\\sc i} $\\lambda$ 6162 \\AA\\ line were also compared with a high resolution spectrum of a very metal-poor star. The conclusions of this work can be summarized as follows: \\begin{itemize} \\item Pressure shifts are shown not to be an important contributor to the asymmetry observed in several solar spectral lines corresponding to neutral elements, but the conclusion cannot be generally extended. \\item The differences in the collisional widths theoretically predicted by Spielfiedel et al. (1991) for the three lines of the Ca\\,{\\sc i} $4s4p-4s5s$ multiplet are reflected in the solar spectrum. \\item The effect of the collisions are negligible when analyzing line asymmetries in very metal-poor stars. \\end{itemize} The comparison made by \\markcite{anstee95}Anstee \\& O'Mara (1995) between their results on line broadening and the solar spectrum give confidence to their analysis. The extension of those calculations to shift cross-sections would be very important. In particular, the application of the detailed molecular-like potential methods to many lines of astrophysical interest. This would allow to analyze the expected effects in different lines, and would help to prevent errors in the studies of velocity fields using line asymmetries. This should be carried out by promoting collaborations between theoretical and experimental physicists and astrophysicists." + }, + "9701/astro-ph9701127_arXiv.txt": { + "abstract": "We present the results of \\asca observations of a heterogenous sample of 15 spiral galaxies. 8 are LINERs or low-luminosity AGN (LLAGN), 5 are starburst galaxies and 2 are normal spiral galaxies. We find that in all cases the \\asca spectra can be described by a canonical model consisting of a power-law with a photon index, $\\Gamma \\sim 1.7-2.0$, plus a soft optically thin emission component with $kT \\sim 0.6-0.8$ keV. The implied element abundances are often sub-solar. The soft component is usually extended and the nuclear, point-like emission is sometimes absorbed by column densities in the range $\\sim 10^{21} - 10^{23} \\rm \\ cm^{-2}$. The relative luminosities of the soft and hard components vary from galaxy to galaxy. For the LINERs, the 2-10 keV luminosity of the hard component is typically $\\sim 10^{40-41} \\rm \\ ergs \\ s^{-1}$ whereas the 0.5-2.0 keV luminosity of the soft component is typically $\\sim 10^{39-40} \\rm \\ ergs \\ s^{-1}$. For starbursts, the 2-10 keV luminosity of is $\\sim 10^{39-40} \\rm \\ ergs \\ s^{-1}$, somewhat lower than the corresponding luminosity of most of the LINERs in our sample. The hard component is similar to the observed X-ray spectra of quasars and also to the {\\it intrinsic} X-ray spectra of classical Seyfert galaxies. Most of the galaxies in our sample exhibit no significant ($\\Delta I/I > 20\\%$) short-term variability (with timescales of a day or less) whereas long-term variability is common. We present a case study of the LINER M81 in detail where there is evidence of large-amplitude ($\\Delta I/I \\sim 70\\%$) variability over several weeks. There is also clear evidence for a broad, complex Fe-K emission line which is compatible with an origin in an accretion disk viewed at $\\sim 40$ degrees. These results suggest a strong connection between classical AGN, LINERs, and starburst galaxies. ", + "introduction": "Optical emission lines are found in the nuclei of three types of galaxies: active galactic nuclei (AGN), low-ionization emission line regions (LINERs; Heckman, 1980) and starburst galaxies. The dominant physical process driving the emission lines in AGN is thought to be accretion onto a supermassive (M $\\sim \\rm \\ 10 ^{5-9} M_{\\sun}$) black-hole, while in starbursts the lines are the result of a current or recent episode of star formation. However, the physical origin of the line emission in LINERs is still unclear (see review by A. Fillipenko in this volume). The two most likely scenarios are AGN-type accretion or starburst-type activity (i.e., shocks from supernovae). X-ray observations provide important clues to distinguish between them. AGN are compact ($\\ll 1$ pc), variable sources in which the X-ray emission has a nonthermal, power-law form. In addition, Seyfert galaxies are characterized by strong Fe K line emission produced by fluorescence in cold matter. On the other hand, nuclear starbursts are frequently extended over kpc scales and are expected to have thermal spectra characterized by coronal X-ray emission lines. Below we shall see that the \\asca (See Tanaka, Inoue \\& Holt 1994) spectra of low-luminosity spiral galaxies are complex, usually requiring a two-component description. Previous X-ray studies have been hampered by the lack of instrument sensitivity, spectral resolution and restricted energy bandpass. In particular, the {\\it hard} X-ray spectra were studied with non-imaging instruments where confusion of more than one bright source in the galaxy was a hinderance. The X-ray satellite \\asca, launched in 1993, is capable of 0.4-10 keV imaging, spectroscopy and temporal analysis. \\asca is the first imaging X-ray satellite sensitive above 4 keV and has better spectral resolution ($\\rm \\Delta E/E \\sim 2\\%$ at 6 keV) than any X-ray imaging mission to date. The half-power radius of the X-ray mirrors is $\\sim 1.'5$ and this is sufficient to obtain individual spectra of sources separated by $\\sim 3-4'$ or more (typical of the bright sources found in galaxies in our sample). Here we present preliminary \\asca results on a sample of fifteen galaxies (some previously published), eight of which are LINERS/LLAGN. When available, some limited analysis was also performed with public \\rosat data. ", + "conclusions": "\\begin{figure}[htbp] \\plotfiddle{lumin_ap_idl.ps} {3.in}{0}{70}{40}{-225}{-50} \\caption{\\small {\\bf a)} 0.5-10 keV luminosities of the galaxy sample. Solid line shows LINERs/LLAGN only. Top axis shows log of number of solar-mass binaries emitting at the Eddington luminosity required to produce the luminosities shown. {\\bf b)} As in a) but with hardness ratios. The dotted lines show the expected hardness ratio from an absorbed power-laws with photon indices of 2.5, 2.0, 1.5 and 1.0 \\normalsize} \\end{figure} \\begin{figure}[htbp] \\plotfiddle{halpha_gamma_ap_idl.ps} {2.65in}{0}{70}{65}{-220}{-250} \\caption{\\small {\\bf a)} Distribution of $\\rm L_X / L_{H\\alpha}$ for AGN observed in the 0.5-4.5 keV band (dashed line) and 2-10 keV band (solid line) from Elvis, Soltan, \\& Keel (1984), with the filled bars showing the \\asca results in the 2-10 keV band for 8 LINERs/LLAGN. {\\bf b)} Distribution of photon indices ($\\Gamma$) observed in the power-law component of LINERS/LLAGNs (dashed line) and the Nandra \\& Pounds (1994) sample of Seyfert 1 AGN (dotted line) (the Nandra \\& Pounds (1994) slopes are from fits including Compton reflection) \\normalsize} \\end{figure} The hard X-ray emission of both starburst and LINER galaxies is usually dominated by a compact nuclear, or near-nuclear source and there are often one to several other off-nuclear X-ray sources in the galaxy which together constitute the bulk of the X-ray luminosity. Figure 10a shows a histogram of the luminosities observed in the entire galaxy sample, with LINERs/LLAGN (NGC 3147, NGC 3998, NGC 4579, NGC 4594, NGC 4258, M51, NGC 3079, M81) plotted separately. LINER/LLAGN are more luminous than the starburst galaxies, in general. In both cases, tens to thousands of solar-mass X-ray binaries radiating at the Eddington limit would be required to produce the observed flux. Also shown (Figure 10b) is a histogram of the hardness ratio (2-10 keV flux / 0.5-2.0 keV flux) from which it can be seen that there is no clear separation between the starbursts and the LINER/LLAGN. Most of the galaxies have a hardness ratio similar that expected from an unabsorbed power-law with $\\Gamma \\sim 1.5-2.5$. An exception is NGC 3628 (traditionally classified as a starburst) which has an exceptionally flat X-ray spectrum with $\\Gamma \\sim 1.2$ (Yaqoob \\etal 1995a). Variability has been observed in both LINER and starburst galaxies on timescales of weeks to years. This implies that in both starbursts and LINERs, a significant amount of the emission is originating from a compact ($\\ll 1$ pc) region, most likely due to a single object such as a LLAGN or X-ray binary (although such a binary would have to have a much larger mass than known X-ray binaries, typically $\\sim 10-1000 \\rm \\ M_{\\sun}$. On the other hand, rapid variability, on timescales of a day or less, is {\\it not} observed. Thus, if LLAGN are simply low-luminosity versions of classical AGN, it is puzzling that rapid variability is most common in objects with 2--10 keV luminosity in the range $\\sim 10^{42} - 10^{43} \\ \\rm erg \\ s^{-1}$ but vanishes below $\\sim 10^{41} \\ \\rm erg \\ s^{-1}$. These results strongly imply a connection between starburst and LINER activity, with the soft component most likely being produced by warm gas with $kT \\sim 0.6-0.8$ keV, possibly from an SNR-heated ISM and starburst-driven winds. In some cases, the hard component of both starburst and LINERs may also be due to starburst activity, possibly resulting from compact supernovae (cSNR). It is possible that some LINERs (and starbursts) may be powered by AGN-type accretion. To test this hypothesis we have plotted histograms of the ratio of X-ray flux to $\\rm H_{\\alpha}$ flux and the photon indices ($\\Gamma$) observed in the LINERs/AGN along with the same quantities for AGN in Figures 11a and 11b, respectively. There is no obvious distinction between the two groups. Unfortunately, our sample is too small to statistically test if the two samples originated from the same parent population, however, Figures 11a and 11b suggest that this is the case. The X-ray spectra of LINERs and starbursts are not very different to those of classical AGN. The hard power-law slopes are similar to those of quasars and the inferred {\\it intrinsic} slopes of Seyfert galaxies. It appears that quasars and the low-luminosity spiral galaxies are devoid of large amounts of matter residing in the nucleus, which is responsible for reprocessing the X-ray continua of the intermediate luminosity Seyfert galaxies (e.g. Nandra \\& Pounds 1994). Even the soft extended thermal emission which is common in many of the LINERs and starbursts has been observed in classical AGN in which the hard power-law is heavily absorbed, allowing the soft component to be detected (e.g. NGC 4151 and Mkn 3; see Figure 4). Another puzzling result is the lack of significant Fe-K line emission in galaxies like NGC 253 and M82 which are thought to be prototypical starburst galaxies in which case the hard X-ray emission would be expected to have a thermal origin. Yet, the lack of Fe-K line emission and the OSSE detection of NGC 253 (Bhattacharya \\etal 1994) suggest a nonthermal origin. Fe-K line emission (in all cases, likely to be due to fluorescence) is clearly detected in only three of the galaxies: NGC 3147 (Ptak \\etal 1996), NGC 4258 (Makishima \\etal 1994) and M81 (Ishisaki \\etal 1996 and \\S 6). In the first two cases the Fe-K line very likely originates in obscuring matter around the nucleus, in close analogy to the mechanism in Seyfert 2 galaxies. In the case of M81 the characteristics of the Fe-K line are very similar to those in Seyfert 1 galaxies and the line may indeed have the same origin; an X-ray illuminated accretion disk. The inferred low accretion rate of M81 may however be problematic. Much work remains to be done in understanding the physical implications of the X-ray emission in low-luminosity spiral galaxies." + }, + "9701/astro-ph9701089_arXiv.txt": { + "abstract": "We have obtained $V$, $R_c$, $I_c$ HRCAM images and TIGER spectrography of the central region of the peanut galaxy NGC~128. The colour images reveal the presence of a red disc tilted by about 26 degres with respect to the major-axis of the galaxy. This tilted disc is made of dust and gas, as revealed by the 2D TIGER map of the ionized gas distribution. The TIGER stellar and gas velocity fields show that the angular momentum vectors of the stellar and gaseous components are reversed. We therefore suggest that the gas orbits belong to the so-called anomalous family, which is evidence for a tumbling triaxial potential (a bar) associated with the peanut morphology. The bar formation has very probably been triggered through the interaction with its nearby companion NGC~127, from which the dissipative component is being accreted along retrograde orbits. ", + "introduction": "N body simulations have shown that strong bars can appear as peanut-shaped bulges when viewed close to edge-on (Combes \\& Sanders 1981, Combes et al. 1990, Pfenniger \\& Friedli 1991). The same bar would be seen nearly round when end-on and box-shaped for an intermediate viewing angle. These morphologies are indeed observed in galaxies, which suggested that boxy and peanut-like bulges could be linked to the presence of bars. Recently, Kuijken \\& Merrifield (1995) have shown that it is possible to detect edge-on bars kinematically by a careful study of the projected velocity distribution. When a tumbling triaxial structure is present, the observed Line Of Sight Velocity Distributions (LOSVDs) should exhibit contrasted gaps corresponding to the transition regions between the main resonances. They indeed detected these expected features in two peanut-shaped galaxies. A larger sample has recently been analysed by Bureau \\& Freeman (1997) who found the signature of a bar in nearly all systems. This technique is powerful to detect bars in edge-on systems. However it requires the presence of a rather extended gas disc\\footnote{Since star orbits can ``cross'' each other, the bar signature is weaker in stellar LOSVDs (see Kuijken \\& Merrifield\\ 1995).}. It may therefore be a difficult task to apply this method systematically on all boxy/peanut bulges. In this Letter, we show that in exceptional circumstances the signature of the bar is even more obvious. This is illustrated with the prototype of peanut-shaped galaxy, namely NGC~128. Although this ``peculiar'' S0 galaxy appears in the Hubble Atlas, it has been only scarcely studied. Bertola \\& Capaccioli (1977) published a combined photometric and spectroscopic analysis of NGC~128, including the derivation of its major-axis velocity profile. It was also included in the sample of early-type galaxies observed by Bertola et al. (1992) who obtained long-slit spectrography along the major-axis, and only commented on the mean angular momenta of the stars and gas concluding that they had the same direction. This is, as will be shown in this Letter, inconsistent with our TIGER data (preliminary results published in Monnet et al. 1995; see also Pagan et al. 1996 and Kuijken et al. 1996). Details on our observations are given in Sect.~\\ref{sec:obs}. The corresponding results are presented in Sect.~\\ref{sec:res}. A brief discussion and some conclusions are drawn in Sect.~\\ref{sec:conc}. ", + "conclusions": "\\label{sec:conc} The regularity of the velocity field of the gas suggests that it has settled onto closed orbits in the centre. In an axisymmetric potential, gas would rapidly fall onto the equatorial plane of the galaxy in a few orbital periods\\footnote{The orbital time is estimated to be $\\sim 26$~Myr at $2\\arcsec$.}. This process seems to be even more efficient if the potential is triaxial and stationary (Colley \\& Sparke 1996). However, tumbling triaxial potentials are known to contain stable families of closed orbits which leaves the plane perpendicular to the rotation axis (see e.g. Magnenat 1982, Mulder \\& Hooimeyer 1984). The main family of retrograde orbits is then the so-called ANomalous Orbits (ANO) which corresponds to the 1:1 vertical resonance. For slow figure rotation $\\Omega_p$ these orbits bifurcates from the E$_z$ family. Above a certain critical value $\\Omega_{p_{crit}}$, the $z$-orbit become complex unstable. The value of $\\Omega_{p_{crit}}$ is significantly lowered as soon as a small central mass concentration is present (see Martinet \\& Pfenniger 1987, and Pfenniger \\& Friedli 1991 for further details). We therefore conclude that the potential of NGC~128 must be triaxial and tumbling. In other words, the prominent peanut in NGC~128 corresponds to a bar viewed nearly edge-on. The maximum of the peanut corresponds to the axisymmetric horizontal and vertical inner Lindblad resonances: $\\Omega_p = \\Omega - \\kappa / 2 = \\Omega - \\nu_z / 2$ (Combes et al. 1990). In a forthcoming paper, we will present a realistic mass model of NGC~128 built using the generalization of the Multi-Gaussian Expansion method (Emsellem 1993). Preliminary results from this model indicate that the pattern speed must be high with $\\Omega_p > 30$~km.s$^{-1}$.kpc$^{-1}$, and therefore suggests that it overcomes the critical boundary $\\Omega_{p_{crit}}$. The morphology of the retrograde gas orbits in such a case (``fast $\\Omega_p$'') are nicely illustrated in Fig.~3 of Friedli \\& Udry (1993): they are circling the long-axis of the bar, and their edge-on projection shows a central tilted disc structure very similar indeed to the one observed in NGC~128. A detailed dynamical model is required to confirm that this interpretation is fully consistent with the observed kinematics. The peanut itself is not perfectly symmetric and exhibits some distortions: this is certainly the result of the interaction with NGC~127. We suggest that the interaction with the small satellite NGC~127 (whose integrated luminosity in the $I_c$ band is $\\sim 7\\%$ of NGC~128's) triggered the formation of a bar in NGC~128 and accelerated its dynamical evolution leading to its peanut-shape. There is a gap between the two discs along the major-axis which closely corresponds to the location of the maximum of the peanut. This is expected if the stars populating the peanut were driven out of the equatorial plane of the galaxy. This double disc structures intriguingly ressembles the ones observed in many S0s (Seifert \\& Scorza 1996). The observed ionized gas in NGC~128 has almost certainly an external origin. Our photometric data show that NGC~127 contains a significant amount of dust (see Fig.~\\ref{fig:hrc}) and therefore very probably some gas. It is however difficult to say whether this has been accreted to form the tilted red disc observed today in NGC~128: it could have been formed during an earlier accretion event. Finally, if NGC~127 is later cannibalized by NGC~128, it is likely that the bar would be destroyed in the process leading to a nearly axisymmetric bulge (Pfenniger 1991)." + }, + "9701/astro-ph9701040_arXiv.txt": { + "abstract": "We analyse the spatial clustering properties of a new catalogue of very rich galaxy clusters selected from the APM Galaxy Survey. These clusters are of comparable richness and space density to Abell Richness Class $\\geq 1$ clusters, but selected using an objective algorithm from a catalogue demonstrably free of artificial inhomogeneities. Evaluation of the two-point correlation function $\\xi_{cc}(r)$ for the full sample and for richer subsamples reveals that the correlation amplitude is consistent with that measured for lower richness APM clusters and X-ray selected clusters. We apply a maxmimum likelihood estimator to find the best fitting slope and amplitude of a power law fit to $\\xi_{cc}(r)$, and to estimate the correlation length $r_{0}$ (the value of $r$ at which $\\xi_{cc}(r)$ is equal to unity). For clusters with a mean space density of $1.6\\times 10^{-6}\\hmpccc$ (equivalent to the space density of Abell Richness $\\geq 2$ clusters), we find $r_{0}=21.3^{+11.1}_{-9.3} \\hmpc$ ($95\\%$ confidence limits). This is consistent with the weak richness dependence of $\\xi_{cc}(r)$ expected in Gaussian models of structure formation. In particular, the amplitude of $\\xi_{cc}(r)$ at all richnesses matches that of $\\xi_{cc}(r)$ for clusters selected in N-Body simulations of a low density Cold Dark Matter model. ", + "introduction": "Rich clusters of galaxies have been used by many authors as tracers of the large-scale structure of the Universe. Most analyses to date have relied on the cluster catalogue of Abell (1958) (and later Abell, Corwin \\& Olowin 1989 (ACO)). Angular clustering statistics for Abell clusters were calculated by Bogart \\& Wagoner (1973) and Hauser \\& Peebles (1973) and more recently, various redshift surveys of Abell clusters have been used to estimate the two point cluster correlation function $\\xi_{cc}(r)$ (eg. Bahcall \\& Soneira 1983, Klypin \\& Kopylov 1983, Postman, Huchra \\& Geller 1992, Peacock \\& West 1992). >From these studies, the two point correlation function for clusters has been found to be consistent in shape with the power law form measured for galaxies, \\begin{equation} \\xi_{cc}(r)=\\left(\\frac{r}{r_{0}}\\right)^{-\\gamma}. \\end{equation} with a similar value of the power-law index $\\gamma\\sim 2$ but with a higher amplitude $r_{0}$. For example, Peacock \\& West (1992) find $r_{0}=21 \\hmpc$ (where H$_{0}=100h \\kms$) for Abell clusters of richness $R \\simgt 1$ whereas $r_{0}$ is around $5 \\hmpc$ for galaxies (see eg. Davis \\& Peebles 1983). Many authors have found, however, that there is much evidence to suggest that the Abell catalogue, selected by eye from unmatched photographic plates, is affected by inhomogeneities in cluster selection which result in articifial clustering (Sutherland 1988; Sutherland \\& Efstathiou 1991; Dekel {\\it et~al.} 1989; Peacock \\& West 1992). New results on the distribution of clusters have been obtained from an automatically selected catalogue based on the APM Galaxy Survey (Dalton {\\it et~al.} 1992, hereafter DEMS92), and from smaller samples of clusters selected from the Edinburgh--Durham Galaxy Catalogue and from the ROSAT X-ray cluster survey (Nichol {\\it et~al.} 1992; Romer {\\it et~al.} 1994) . The amplitude of $\\xi_{cc}$ measured from these studies is generally lower than for the Abell samples, so that $13 \\hmpc \\simlt r_{0} \\simlt 16 \\hmpc$. However, it has been argued that the clustering seen in the automated surveys is dominated by poor clusters, and that the results may be compatible with the higher values of $r_{0}$ measured for $R \\simgt 1$ Abell clusters, provided that there is a strong dependence of the correlation length on cluster richness. Bahcall \\& West(1992) and Bahcall \\& Cen (1992) argue that the Abell data are consistent with a linear relation between $r_{0}$ and mean intercluster separation $d_{c}$ ($d_{c}=n_{c}^{-1/3}$ where $n_{c}$ is the mean space density) so that \\begin{equation} r_{0}=0.4 d_{c}. \\end{equation} The evidence for this scaling relation, especially at high values of $d_{c}$ comes exclusively from estimates of the correlation functions of rich Abell clusters (eg. Peacock \\& West 1992). The validity of equation (2) thus depends critically on the uniformity of the Abell catalogue, particularly at richnesses $R \\simgt 1$. The main aim of this paper is to test the scaling relation (2) using an independent sample of rich clusters of galaxies selected from the APM galaxy survey. Croft \\& Efstathiou (1994) have shown that the amplitude of the cluster correlation function is is predicted to vary only weakly with cluster space-density for a range of Cold Dark Matter (CDM) models. The existing data for APM clusters (Dalton {\\it et~al.} 1994a) are in good agreement with these predictions, but as the clusters are of relatively low richness and hence low $d_{c}$ they are also consistent with the relationship given in Equation 2. In the study presented here we use a new extension of the APM cluster survey to test the behaviour of $\\xi_{cc}$ for richer clusters. The layout of this paper is as follows. We describe the cluster sample and its relationship to the samples of Dalton {\\it et~al.}(1994a) and Dalton {\\it et~al.} (1994b) in Section 2. In Section 3 we present the correlation function for the new cluster sample and for various subsamples . We use a maximum likelihood estimator to fit a power law to the correlation function and investigate how the fitted parameters change with cluster richness. In Section 4 we compare our results with other data samples and with N-body simulations of cosmological models. We summarise our findings in Section 5. ", + "conclusions": "\\subsection{Comparison with results for other cluster catalogues.} A plot of $r_{0}$ versus $d_{c}$ for various observational samples of clusters including the APM samples (taken from Table 2 above and DEMS92 and labelled APM and APM92 respectively). is shown in Fig 4. For the maximum likelihood points, the errors are $1 \\sigma $ for marginalisation of $r_{0}$ over all values of $\\gamma$. The APM92 points are for $\\gamma$ constrained to be 2.0. The points labelled `Abell' indicate the results for Abell $R \\geq 0$ , Abell $R \\geq 1$ and Abell $R \\geq2$ clusters derived by Peacock \\& West (1992). The point labelled EDCC is the result for 79 clusters from the Edinburgh-Durham Cluster Catalogue of Lumsden {\\it et~al.} (1992) estimated by Nichol {\\it et~al.} (1992). The error bar size was estimated using bootstrap resamplings. The same is true of the error on the point labelled X-Abell, which was estimated by Nichol, Briel \\& Henry (1994), from 67 clusters in the redshift sample of Huchra {\\it et~al.} (1990) which also have X-ray luminosities $\\geq 10^{43}$ erg s$^{-1}$. The point labelled ROSAT shows $r_{0}$ for $\\xi_{cc}(r)$ measured from a redshift survey of an X-ray flux limited sample of clusters (Romer {\\it et~al.} 1994). The X-ray flux for both these last samples was measured using the ROSAT satellite. \\begin{figure} \\centering \\vspace{7.5cm} \\special{psfile=\"r0vsdobs.ps\" angle=-90 vscale=45 hscale=45 voffset=260 hoffset=-45} \\caption[junk]{\\label{junk2} The quantity $r_{0}$ (the correlation length) plotted against cluster space density for a number of observed cluster samples (see text). Error bars represent the 1 $\\sigma$ error on the mean. The solid line shows the relation r$_{0}= 0.4 d_{c}$ of Bahcall \\& West (1992).} \\end{figure} It can be seen that most of the data points are for cluster samples with $d_{c}$ in the range $30-55\\hmpc$, and that in this range, the results for the X-ray samples and automated galaxy surveys are in agreement with one another, and lower than those for Abell clusters. As has been detailed previously, this can be understood as being due to non-uniformities in the Abell catalogue which artificially boosts the amplitude of clustering. Over this small range in cluster space density, for which the errors are comparatively small, there is not much evidence for any trend of $r_{0}$ with $d_{c}$ and hence cluster richness. Part of the reason for the work in this paper was to find out whether this is also true at higher richnesses and lower space densities. The solid line in the plot corresponds to the scaling relation $r_{0}=0.4d_{c}$ proposed by Bahcall \\& West (1992) as a fit to the correlation functions of the Abell sample. As can be seen from the plot, the motivation for assuming this fit at high values of $d_{c}$ was provided by the results for Abell $R \\geq 2$ clusters (a sample of 42 clusters was used to calculate this data point -- see Peacock \\& West 1992). Now that we have a sample of very rich clusters taken from a catalogue which is demonstrably free of artificial inhomogeneities, we are in the position to test equation (2) using the APM data alone. The three APM points on the right of the plot are for ${\\cal{R}}\\geq 90, {\\cal{R}}\\geq 100 $ and ${\\cal{R}}\\geq 110$ clusters, which have space densities comparable to that of the Abell $R \\geq 2$ clusters. If the error bars are taken at face value, then the relation would appear to be ruled out at the $\\sim 2 \\sigma$ level. However, as we have seen from Table 1, the error bars could be underestimates by a factor of $\\sim1.1 - 2.1$. Also, the space densities of clusters used to derive $d_{c}$ values are not precise estimates because of the difficulties involved in estimating the completeness of richness limited cluster catalogues (see Efstathiou {\\it et~al.} 1992). That said, we believe that these data points are more reliable than those for the Abell $R \\geq 2$ clusters. Table 2.1 also shows us that the error bars for the richest sub-samples are likely to be the most accurate. In summary, the APM points are consistent with a weak dependence of clustering on richness. We find no evidence that equation (2) applies to rich clusters of galaxies, with important implications for theories of structure formation as described in the next section. \\subsection{Comparison with model predictions.} Croft \\& Efstathiou (1994) examined the behaviour of $r_{0}$ with $d_{c}$ expected in several popular cosmological scenarios (see also Bahcall \\& Cen 1992, Mann, Heavens \\& Peacock 1993). The box size ($300 \\hmpc$) of the dissipationless N-body simulations used in that study, meant that the predictions did not extend to the large values of $d_{c}$ needed to make comparisons with our new rich cluster sample. We have therefore run a set of simulations (using the same particle-particle particle-mesh N-body code) with box size $600 \\hmpc$ and $4 \\times 10^{6}$ particles. These simulations are the same as those used in Croft \\& Efstathiou (1995). The models we shall consider are the Standard CDM model (SCDM has $\\Gamma= \\Omega h=0.5$ and $\\Omega=1$) and the spatially flat Low density CDM model (LCDM has $\\Gamma=0.2$, $\\Omega=0.2$ and $\\Omega_{\\Lambda}=0.8$). Both models are normalised to be compatible with the first year COBE anisotropies (Wright {\\it et~al.} 1994) so that $\\sigma_{8}=1.0$ for both models, where $\\sigma_{8}$ is the rms amplitude of linear fluctuations in $8\\hmpc$ spheres. The results are insenstive to the precise value of $\\sigma_{8}$ . We use the same techniques as in Croft \\& Efstathiou (1994) and Dalton {\\it et~al.} (1994a) to find clusters in the simulations. This involves finding cluster centres in real space with a percolation algorithm and then ordering clusters by the mass contained within a certain radius, in this case $0.5 \\hmpc$. We then calculate $r_{0}$ for clusters with different lower mass limits, with the results shown in Figure 5. We have chosen to calculate the correlation functions in redshift space, for more accurate comparison with the observations. The values of $r_{0}$ which we present below are $\\sim 1 \\hmpc$ larger than the values estimated in real space. The correlation functions for the LCDM model are shown in Figure 5, together with the APM points (estimated using Equation 4). The space densities of the simulated clusters were selected to be close to those for the three subsamples of rich APM clusters plotted. The curves plotted are the averages of results for 10 simulations of LCDM. We can see that the APM results are compatible with LCDM model. We can also see that the clustering strength of LCDM clusters increases only a small amount as the richness bound is increased. \\begin{figure} \\centering \\vspace{8.7cm} \\special{psfile=\"xsirichlcdm.ps\" angle=-90 vscale=50 hscale=50 voffset=290 hoffset=-70} \\caption[junk]{\\label{junk2} The two-point correlation function of simulated clusters in the LCDM model in redshift space for subsamples with three different mean separations. The APM results from Figure 2 (solid symbols, computed using the estimator of Equation 4) are also shown. } \\end{figure} In order to see how the clustering results are affected by the mask and selection function, we have plotted the results for the mock APM cluster catalogues constructed from 10 LCDM simulations and described in Section 3.1. The results are shown in Figure 6 in the form of a scatter plot. In each panel we plot $r_{0}$ against $\\gamma$ (both measured from the maximum likelihood technique). We show results calculated from clusters with the same $d_{c}$ values as those in Figure 5. We also plot points for the APM results for equivalent richess clusters. In each of the panels we can see that the APM results are not extreme outliers and it looks plausible that they could have been drawn from the same distribution as the LCDM points. A line denoting the relationship of Equation 2 is drawn on each panel. From this we can conclude that in an LCDM Universe we would have a $\\sim 10 \\%$ chance for each richness cut of measuring a value of $r_{0}$ which fits this relationship. \\begin{figure} \\centering \\vspace{14.7cm} \\special{psfile=\"r0gammscat.ps\" angle=-90 vscale=80 hscale=80 voffset=475 hoffset=-190} \\caption[junk]{\\label{junk2} Values of $r_{0}$ and $\\gamma$ measured from mock catalogues constructed from 10 simulations of an LCDM universe (see text). Plotted are results for clusters with three different mean separations. The corresponding APM results from Figure 3. are also shown. } \\end{figure} \\begin{figure} \\centering \\vspace{7.5cm} \\special{psfile=\"r0vsdrich.ps\" angle=-90 vscale=45 hscale=45 voffset=260 hoffset=-45} \\caption[junk]{\\label{junk2} A comparison of the richness dependence of APM cluster correlations (filled circles) with the corresponding predictions for a low density CDM Model (dashed line) and Standard CDM (dot-dashed) line. The theoretical predictions have been calculated in redshift space. Error bars represent the 1 $\\sigma$ error on the mean. The solid line shows the relation r$_{0}= 0.4 d_{c}$ of Bahcall \\& West (1992).} \\end{figure} In Figure 7 we plot $r_{0}$ measured from the correlation functions of the LCDM and SCDM clusters against $d_{c}$. The error bars on the simulation points were calculated from the $1 \\sigma$ error on the mean taken from 3 simulations of each model. We therefore have 2.4 times as many clusters of any given space density as in the ensembles of Croft \\& Efstathiou (1994). We also plot the values of $r_{0}$ calculated using the maximum likelihood method in this paper. The plots shows the very weak trend of clustering strength with cluster richness continuing for both models at least up to $r=80 \\hmpc$. The APM points are consistent with the LCDM model, but not with SCDM. We note here that a simulation with a different amplitude of clustering in the underlying mass could have an $r_{0}$ which differs by as much as $1-3 \\hmpc$ as could clusters which are selected using a different method. These variations are not expected to be large enough to affect our conclusions (Croft \\& Efstathiou 1994, Eke {\\it et~al.} 1996, Mo, Jing \\& White 1996). It is encouraging that models with $\\Gamma \\approx 0.2$, which were introduced to explain clustering in the galaxy distribution (see eg. Efstathiou, Sutherland \\& Maddox 1990) are also able to match well the clustering of rare and extreme objects such as the rich galaxy clusters considered here. We also expect other Gaussian models with similar power spectra such as a Mixed Dark Matter (MDM) universe dominated by CDM and with an additional component of massive neutrinos (see eg. Klypin {\\it et~al.} 1993) to be compatible with our APM results at high richnesses, as they are at low richnesses (Dalton {\\it et~al.} 1994a). >From Figure 7 we can also see that whilst models such as low density CDM provide a good fit to the clustering behaviour of rich APM clusters they are completely incompatible with the scaling relation derived from considering rich Abell clusters. Our data exclude such a strong scaling relation and remove the need to resort to non-Gaussian models for the formation of large-scale structure." + }, + "9701/astro-ph9701106_arXiv.txt": { + "abstract": " ", + "introduction": "Protostellar discs appear to be common around young stars. Furthermore recent studies show that almost all young stars associated with low mass star forming regions are in multiple systems (Mathieu, 1994 and references therein). Typical orbital separations are around 30 astronomical units (Leinert et al. 1993) which is smaller than the characteristic disc size observed in these systems (Edwards et al. 1987). It is therefore expected that circumstellar discs will be subject to strong tidal effects due to the influence of binary companions. \\noindent The tidal effect of an orbiting body on a differentially rotating disc has been well studied in the context of planetary rings (Goldreich and Tremaine, 1981), planetary formation, and generally interacting binary stars (see Lin and Papaloizou, 1993 and references therein). In these studies, the disc and orbit are usually taken to be coplanar (see Artymowicz and Lubow, 1994). However, there are observational indications that discs and stellar orbits may not always be coplanar ( see for example Corporon, Lagrange and Beust, 1996 and Bibo, The and Dawanas, 1992.) \\noindent In addition, reprocessing of radiation from the central star by a warped non coplanar disc has been suggested in order to account for the high spectral index of some T~Tauri stars (Terquem and Bertout 1993, 1996). \\noindent A dynamical study of the tidal interactions of a non coplanar disc is of interest not only in the above contexts, but also in relation to the possible existence of precessing discs which may define the axes for observed jets which apparently precess (Bally and Devine, 1994). \\noindent Various studies of the evolution of warped discs have been undertaken assuming that the forces producing the warping were small so that linear perturbation theory could be used ( Papaloizou and Pringle, 1983, Papaloizou and Lin, 1995 and Papaloizou and Terquem, 1995). The results suggested that the disc would precess approximately as a rigid body if the sound crossing time was smaller than the differential precession frequency. We describe here some recent non linear simulations of discs which are not coplanar with the binary orbits using a Smoothed Particle Hydrodynamics (SPH) code originally developed by Nelson and Papaloizou (1993, 1994). We study the conditions under which warped precesing discs may survive in close binary systems and the truncation of the disc size through tidal effects when the disc and binary orbit are not coplanar. The simulations indicate that the phenomenon of tidal truncation is only marginally affected by lack of coplanarity. Also our model discs were able to survive in a tidally truncated condition while warped and undergoing rigid body precession provided that the Mach number in the disc was not too large. The inclination of the disc was found to evolve on a long timescale likely to be the viscous timescale, as was indicated by the linear calculations of Papaloizou and Terquem (1995). ", + "conclusions": "We have described nonlinear simulations of an accretion disc in a close binary system when the disc midplane is not necessarily coplanar with the plane of the binary orbit. For our constant viscosity SPH models we found the tidal truncation phenomenon to be only marginally affected by non coplanarity. We found that modestly warped and thin discs undergoing near rigid body precession may survive in close binary systems. However, extremely thin discs may be severely disrupted by differential precession depending on the magnitude of the characteristic Mach number, ${\\cal M}.$ The crossover between obtaining a warped, but coherent disc structure, and disc disruption occurs for a value of the Mach number ${\\cal M} \\sim 30.$ We also found that the inclination evolved on a long timescale, likely to be the viscous timescale, as indicated by the linear calculations of Papaloizou and Terquem (1995). \\noindent A class of models formulated to explain the generation of jets in young stellar objects assumes that a wind flows outwards from the disc surface. This is then accelerated and collimated by the action of a magnetic field (see K\\\"onigl and Ruden, 1993 and references therein). It is reasonable to assume that a precessing disc may lead to the excitation of a precessing jet. The precession period obtained from our calculations with a mass ratio of unity is about 500 units. When scaled to a disc of radius 50 AU, surrounding a star of $1 {\\rm M}_{\\odot}$, the unit of time is $\\Omega(R)^{-1} \\simeq 56$~$yr$, leading to a precession period of $3.10^4 yr$. \\noindent Bally and Devine (1994) suggest that the jet which seems to be excited by the young stellar object HH34* in the L1641 molecular cloud in Orion precesses with a period of approximately $10^4$~$yr$. This period is consistent with the source being a binary with parameters similar to those we have used in our simulations with a separation on the order of a few hundred astronomical units. \\noindent Some of the results presented here demonstrate how a warped disc can present a large surface area for intercepting the primary star's radiation. The effect that the consequent reprocessing of the stellar radiation field can have on the emitted spectral energy distribution has been investigated by Terquem and Bertout (1993, 1996). They find that it may account for the high spectral index of some T~Tauri stars or even for the spectral energy distribution of some class~0/I sources. The Model~3 simulation indicates that the required strongly warped disc could be physically realisable. \\noindent Finally, there is evidence from the light curves of X--ray binaries such as Hercules X-1 and SS433, that their associated accretion discs may be precessing in the tidal field of the binary companion ( Schandl and Meyer, 1994). Larwood et al (1996) have demonstrated that the disc precession periods seen in simulations are in reasonable agreement with those that are inferred observationally ( Petterson, 1975, Gerend and Boynton, 1976, Margon, 1984). \\vspace{.2cm}\\noindent{\\it Acknowledgement:} This work was supported by PPARC grant GR/H/09454, JDL is supported by a PPARC studentship." + }, + "9701/astro-ph9701218_arXiv.txt": { + "abstract": "We have analysed HST {\\tt WFPC2} F555W and F814W (i.e., $V$ and $I$) images for fifteen elliptical galaxies with kinematically-distinct cores. For each of them we have derived surface brightness and isophotal parameter profiles in the two bands, color maps, and radial profiles in $V-I$. Most galaxies show patchy dust absorption close to their nuclei. However, there are generally no indications of homogeneous, diffuse dust components close to the nuclei. The nuclear colors in the unobscured regions are most likely representative of the central stellar populations. We have detected photometric evidence for faint stellar disks, on scales of a few tens to a few arcseconds, in seven galaxies, namely NGC 1427, 1439, 1700, 4365, 4406, 4494 and 5322. In NGC 1700, the isophotes are slightly boxy at the scale of the counter-rotating component, and disky at larger radii. We find no difference in $V-I$ color greater than 0.02 mag between these disks and the surrounding galactic regions. Hence the stellar populations in the kinematically distinct cores are not strongly deviant from the population of the main body. Specifically, there is no evidence for a dominating population of blue, very metal weak stars, as predicted by some of the formation scenarios. This argues against models in which small galaxies fall in and survive in the nuclei, unless super massive black holes are present. These would in fact disrupt the accreted small systems. For one galaxy, NGC 4365, the innermost region is bluer than the surrounding regions. This area extends to $\\sim 15$pc, and contains a luminosity of $\\sim 2.5 \\times 10^6$ L$_\\odot$. If interpreted as a stellar population effect, an age difference of $\\sim$ 3-4 Gyrs, or an $[Fe/H]$ variation of about 0.2 dex, is derived. The nuclear intensity profiles show a large variety: some galaxies have steep cusp profiles, others have shallow cusps and a ``break radius''. The nuclear cusps of galaxies with kinematically-distinct cores follow the same trends as the nuclei of normal galaxies. We have not been able to identify a unique, qualifying feature in the {\\tt WFPC2} images which distinguish the galaxies with kinematically distinct cores from the kinematically normal cores. It is possible that statistical differences exist: possibly, the kinematically distinct cores have a higher fraction of nuclear disks. The similarity of both types of cores puts strong constraints on the formation scenarios. Simulations of galaxy mergers, with the inclusion of star formation and nuclear black holes, are needed to resolve the question how these structures may have formed. Spectra with high spatial resolution are needed to study the nuclear structure of the distinct component in detail. ", + "introduction": "The study of nearby (kinematically-normal) elliptical galaxies performed with the {\\tt WF/PC} camera aboard the {\\it Hubble Space Telescope} (HST) has revealed the presence of central disks, dust clouds, and nuclear components on scale lengths of a few tens of parsecs (e.g., Jaffe et al. 1994; van den Bosch et al. 1994; Lauer et al. 1995, hereafter L95; van Dokkum \\& Franx 1995). In all galaxies, the surface brightness profile $I(r)$ for $r \\rightarrow 0.1''$ can be approximated as a power--law $r^{-\\gamma}$, with $0\\lta \\gamma \\lta 1$. Galaxies with steep cusps are almost always small, fast--rotating objects, while in contrast massive, pressure--supported ones have shallow cusps, typically with $\\gamma \\lta 0.5$ (e.g., L95; Kormendy et al. 1995, hereafter K95). It is not clear whether this reflects a real dichotomy in the physical properties and formation processes, as suggested by L95 and K95. In any case, nuclear properties on scales of a few tens of parsecs are correlated with properties observed on global scales, and it is legitimate to ask whether this is the outcome of selection effects acting at formation, or whether it has resulted from subsequent galactic evolution. Galaxies with kinematically-distinct cores are very interesting physical laboratories to test ideas of galaxy formation and evolution. Their distinct angular momenta suggest that these cores may be the relics of interactions or mergers, and thus may provide a powerful diagnostic to estimate the relevance of interactions in shaping galactic properties. Kormendy (1984) argued that the core of NGC 5813 is dominated by the remnant of a small galaxy that fell into a big elliptical. This scenario predicts core/global differences that should be detectable not only in the kinematics, but also in the stellar populations and overall structure. Another possible scenario for kinematically-distinct cores implies gas-rich mergers of two galaxies, where the gas is converted into stars in the central regions, and subsequent star formation produces a disk which retains the orbital signature of the original gas. Still another possibility is that formation of kinematically-distinct cores is a rather normal aspect of a hierarchical formation scenario for elliptical galaxies, where galaxies are built up from pre-existing clumps (possibly different from current galaxies). The last two scenarios do not necessarily predict any measurable difference in the stellar population of the distinct core. Forbes, Franx \\& Illingworth (1995, hereafter F95) investigated HST {\\tt WF/PC} F555W images for eight early-type galaxies with kinematically-distinct cores. These authors found clear indications of small-scale dust in six out of the eight galaxies, and argued that the dust might be directly linked to the existence of nuclear radio activity in these galaxies. They also showed that the nuclear surface brightness profiles in this class of galaxies are similar to those in kinematically-normal galaxies. On this basis, they argued against the plausibility of the small galaxy accretion scenario, since this low-mass object should dominate the light in the core region and produce differences with respect to normal galaxies (that are not observed). F95 searched for the nuclear disks expected in the dissipative core formation scenario (and suggested by spectroscopic analysis of several distinct cores, see e.g., Franx \\& Illingworth 1988). The search lead to one detection in NGC 4365 (previously reported by Forbes 1994), and to a couple of other candidates. In the other galaxies, no conclusions could be drawn due to the strong dust absorption. Additional information is clearly needed to resolve the enigma of the formation of kinematically-distinct cores. Multi-color, high-resolution imaging can provide a valuable key for understanding the origin of these peculiar stellar systems. In particular, within the limits imposed by their degeneracy between age and metallicity, {\\it (i)} color gradients between the core region and the rest of the galaxy could be used to estimate metallicity/age differences in the stellar populations; {\\it (ii)} color information could help to detect nuclear disks, and indicate whether such disks were formed at the same time as the majority of stars (or more recently, as a bluer color might indicate); {\\it (iii)} colors of the central cusps could provide insight into their formation, discriminating between accretion of low-mass stellar systems (possibly leading to bluer cusps) and stars formed in situ (possibly leading to red cusps). Observationally, colors and color gradients at subarcsecond scales are practically unexplored for any kind of galaxy. The only other example to date is that of Crane et al. (1993), who measured gradients for 12 galaxies using pre-refurbishment FOC data, and concluded that color variations are either very small or absent in galactic nuclei. However, these measurements are limited by low S/N and the uncertain {\\tt pre-COSTAR} PSF. In this paper we extend the investigation started by F95 by studing {\\tt WFPC2} images in two bands, namely F555W and F814W (i.e., similar to Johnson $V$ and Cousins $I$, respectively), for 15 galaxies with kinematically-distinct cores. The new sample contains the eight galaxies of F95, plus seven additional galaxies that also show peculiar core kinematics. Eleven of the galaxies are luminous objects with absolute $V$ mag $M_V\\le$-20 + 5 log $h$ ($h=H_o/100$ km s$^{-1}$ Mpc$^{-1}$ throughout this paper). The sample, the observations and the basic photometric analysis are described in Section 2. The $V$ and $I$ surface brightness profiles and the parameters defining the core morphology are presented in Section 3, together with the $V-I$ color profiles, the color maps, and the color gradients. In Section 4 we present the adopted parametrization of the surface brightness profiles, and derive the global estimator of the nuclear cusp-steepness which we will use in our discussion. In Section 5 we discuss the main results of this study, and summarize them in (the concluding) Section 6. In four appendices, we describe {\\it (i)} the procedure used to minimize the effects of nuclear patchy dust on the derivation of the surface brightness profiles, {\\it (ii)} the comparison of our {\\tt WFPC2} measurements with {\\tt WF/PC} profiles already published for some galaxies of our sample; {\\it (iii)} the deprojections of the surface density profiles, performed in order to obtain an estimate of the physical luminosity density and of its nuclear radial variation; and {\\it (iv)} the derivation of the nuclear properties of the L95 galaxies, performed so as to compare kinematically-peculiar with normal galaxies. The properties of the globular cluster systems in our sample are derived and discussed by Forbes et al. (1996). ", + "conclusions": "In this paper we have presented an analysis of {\\tt WFPC2} F555W and F814W images for fifteen galaxies with kinematically-distinct cores. The color data were used to correct for the effects dust. We have derived surface brightness and isophotal parameter profiles in the two bands, and color maps and radial profiles in $V-I$. The surface brightness profiles were parametrized by using an average logarithmic slope inside 10-50 pc. This was then used to investigate the nuclear stellar properties of the galaxy cores. A similar parametrization was used for the nuclear stellar morphologies for the kinematically-normal galaxies of L95, so as to compare the results obtained for our sample with a reference sample. We deprojected the surface brightness profiles, and investigated the relationship between projected quantities and physical stellar luminosity densities and cusp slopes. Our main results are: \\begin{itemize} \\item There is no unique photometric parameter that distinguishes the galaxies with kinematically-distinct cores. They show a large range in properties on the scale of the kinematically-distinct cores, like their normal equivalents, namely shallow to steep cusps, and small-scale disks and boxiness. Many of the galaxies show dust, often patchy and chaotic, and sometimes distributed in a ring or disk. \\item With some possible exceptions, no obvious evidence is found for homogeneous, diffuse dust {\\it confined} to the galactic nuclei. Thus, the nuclear colors are most likely representative of the stellar population properties. \\item In one galaxy, namely NGC 4365, the innermost region (15pc in size, for which the total luminosity is $\\sim 2.5 \\times 10^6$ L$_\\odot$) is bluer by about 0.05 mag than the surrounding regions. If this feature is interpreted as due to a variation in the stellar population, synthetic population models suggest an age younger by $\\sim$ 3-4 Gyrs (or an $[Fe/H]$ variation of about -0.2 dex) between the feature itself and the adjacent galactic regions. \\item The radial logarithmic slopes of the nuclear stellar profiles are found to cover a large range of values, from almost flat to very steep (the slopes $\\gamma$ range from $\\sim$ 0.1 to 0.8). Some galaxies show a clear ``break radius'', whereas other don't. The nuclear cusp slopes of kinematically-distinct core galaxies occupy the same region of parameter space as that occupied by kinematically-normal galaxies. Independent of the core kinematical properties, slowly-rotating, anisotropic, high metallicity galaxies preferentially have shallow cusps, while fast-rotating, low-metallicity galaxies tend to have steep cusps. Despite the bimodal distribution of the cusp slopes, the trends of the cusp slope with global galactic properties might be smooth and continuous, rather than appearing as a distinct dichotomy, as indicated by e.g., K95. \\item Photometric evidence is found for faint stellar disks in seven galaxies, namely NGC 1427, 1439, 1700, 4365, 4406, 4494 and 5322. However, no difference in $V-I$ color ($<0.02$ mag) between these disks and the surrounding galactic regions is present. These differences are much smaller than the gradients are larger radii, or the color differences between galaxies (0.15 mag). The nuclear populations are either very homogeneous, or finely-tuned conspiracies between age and metallicity variations exist. High resolution spectroscopy with HST is needed to test for the existence of variations in the metal absorption lines. \\end{itemize} These data have added new constraints on aspects of the formation process of elliptical and early-type galaxies. While dust has certainly complicated the interpretation of the data on very small scales, techniques that minimize its effect can be used and allow one to extract reliable quantitative information on structures, surface brightness and luminosity density distributions, as well as on the nature of the stellar population in the cores from colors. Comparison of the nuclear with the global properties suggests that kinematically-distinct cores might be the normal outcome of regular hierarchical merging at early epochs, and not indicative of any unusual late event in the galaxies evolutionary history. HST spectroscopy will be extremely valuable to study the nuclear structure of the kinematically distinct components. It remains remarkable how these components stand out in the kinematics, and how difficult they are to find in imaging. HST spectroscopy may be able to resolve many of the unanswered questions of this paper." + }, + "9701/astro-ph9701168_arXiv.txt": { + "abstract": "The predictions of the two-phase accretion disc-corona models for active galactic nuclei are compared with observations. We discuss the possibility to use X-ray spectral slopes, equivalent widths of the iron line, and the observed flux-spectral index correlation as diagnostics of the X-$\\gamma$-ray source compactness and geometry as well as of the cold disc temperature. As an example of the application of the modelling tools, we use XSPEC to fit the broad-band data of Seyfert 1 galaxy, IC4329A, with a theoretical spectrum from a hemisphere-corona. \\vspace{5pt} \\\\ Keywords: accretion, accretion discs; galaxies: Seyfert; gamma rays: theory; X-rays: galaxies. ", + "introduction": "Observations of Seyfert 1 galaxies by {\\em Ginga} and earlier by {\\em HEAO 1} showed that their X-ray spectra consist of at least two distinct components: an intrinsic power-law with energy index, $\\a218=$0.95$\\pm$0.15, in 2 - 18 keV range and a Compton reflection bump (\\cite{mu93}, \\cite{np94}, \\cite{w95}). Recent {\\em ASCA} observations also show these features (\\cite{n96}). {\\em OSSE} observations show a much steeper spectrum with $\\alpha\\approx$ 1.5 (\\cite{joh96}), and {\\em COMPTEL} has not detected Seyfert galaxies at all (\\cite{mai95}). Averaged broad-band spectrum for a sample of Seyfert 1 galaxies which was obtained using non-simultaneous {\\em Ginga} and OSSE data shows a high-energy cutoff of the intrinsic power-law at $\\Ec=$600$^{+800}_{-300}$ keV (\\cite{z95}). A similar cutoff energy is obtained by \\cite*{gon96} using the data from {\\em EXOSAT} and OSSE. {\\em ROSAT, Ginga} and OSSE data of the Seyfert 1 galaxy, IC4329A, also require a cutoff in the intrinsic spectrum (\\cite{ma95}). The brightest Seyfert galaxy in $\\gamma$-rays, NGC 4151, shows a cutoff at $\\sim$ 150 keV, but is probably a freak object requiring special considerations (\\cite{p96}, but see \\cite{zdz97}). Non-thermal pair models intensively studied in the mid 80s (\\cite{sve86}, for review see \\cite{sve94}) predicted the spectral index of the intrinsic component $\\alphint\\approx 0.9-1.0$ and a prominent pair annihilation line. Absence of the annihilation line in the observed spectra rules out pure non-thermal models. Attention is now focused on thermal models. A definite answer on the question about the relative fraction of the thermal and non-thermal plasma cannot be obtained without high quality data above 200 keV. \\begin{figure}[h] \\begin{center} \\leavevmode \\epsfysize=5.0cm \\epsfbox[94 380 478 700]{fig1.ps} \\end{center} \\caption{\\em An example of the deconvolution of the emergent spectrum from a disc-corona system into different scattering orders. Solid curve represents the total emergent spectrum. Dotted curves give the contribution of the unscattered radiation which consists of three components: black-body radiation with a maximum at $\\epsilon\\equiv h\\nu/m_ec^2\\approx$3$\\cdot$10$^{-4}$, a Compton reflection component with a maximum at $\\epsilon\\approx$0.1, and a broad annihilation line at $\\epsilon\\approx$2. Dashed curve gives the contribution of the single scattered radiation (notice that the broad bump with a maximum at $\\epsilon\\approx$0.2 is the first scattering order of the Compton reflection component). All other scattering orders are represented by dash-dotted curves. } \\label{fig:decon} \\end{figure} There is a consensus that the X/$\\gamma$-ray spectrum of Seyferts is produced by Comptonization in hot plasmas of soft radiation from the UV source. The exact geometry of both phases is unknown. Presently, the most commonly used scenario is the two-phase disc-corona model (e.g., \\cite{hm91}, 1993) in which a hot X-ray emitting corona is located above the cold UV-emitting disc of the canonical black hole model for AGNs. The power-law X-ray spectrum is generated by thermal Comptonization of the soft UV-radiation. About half of the X-rays enters and is reprocessed by the cold disc, emerging mostly as black body disc radiation in the UV (some fraction is reflected producing a Compton reflection component). \\cite*{hm91} emphasized the coupling between the two phases due to the reprocessing, as the soft disc photons influence the cooling of the corona. They showed that nearly all power must be dissipated in the corona in order to have $\\alphint \\sim$ 0.9-1. A consequence of this is that the soft disc luminosity, $\\Ls$, is of the same order as the hard X-ray luminosity, $\\Lh$. \\begin{figure}[htbp] \\begin{center} \\leavevmode \\epsfxsize=6.0cm \\epsfbox[164 380 418 720]{fig2.ps} \\end{center} \\caption{\\em Upper panel: spectra emerging in different directions from a slab-corona. Solid, dotted, and dashed curves represent the flux at viewing cosine angle $\\mu$=0.9, 0.5 and 0.1, respectively. The temperature of the black-body radiation from the accretion disc is fixed at $\\Tbb$=50 eV. The electron temperature of the corona increases from the top of the figure to the bottom: $\\Theta\\equiv k\\Te/mc^2$= 0.16, 0.2, 0.25, 0.32, 0.4, 0.5, 0.63, 0.79, 1.0. Notice the increase of the iron line equivalent width with increasing temperature due to anisotropic scattering effects (see, e.g., Haardt $\\&$ Maraschi 1993, Stern et al. 1995b). The iron line profile has a triangular shape due to the energy resolution employed here. Lower panel: spectra from a cylinder-corona with a height-to-radius ratio 2. } \\label{fig:slab} \\end{figure} Observations show that $\\Ls$ may be several times larger than $\\Lh$, in contradiction to the prediction of the uniform two-phase disc-corona model. This led \\cite*{hmg94} to propose a patchy disc-corona model, where the corona consists of several localized active regions on the surface of the disc. Internal disc dissipation results in UV-radiation that leaves the disc without entering the active regions, thus enhancing the observed $\\Ls$ relative to $\\Lh$. The first exact radiative transfer/Comptonization calculations in a disc-corona system accounting for energy and pair balance as well as reprocessing by the cold disc (including angular anisotropy) were reported in \\cite{s95b}. Two methods were used. The first method is based on the Non-Linear Monte-Carlo method by Stern (see detailed description in \\cite{s95a}). The second method is the iterative scattering method (ISM) where the radiative transfer is exactly solved for each scattering order separately (\\cite{ps96}). The results of both codes are in excellent agreement. In this paper we report the results using the second method. A typical spectrum emerging from the disc-corona system and its different components is shown in Figure~\\ref{fig:decon}. ", + "conclusions": "Several diagnostics from physical modelling of observed spectra are possible. (1) As anisotropic effects are very important, i.e. as spectral shapes depend strongly on viewing angle, it will be possible to set constraints on the viewing angle within the framework of the two-phase pair corona model, once high quality spectra become available. (2) The spectra also depend on the geometry of the coronal regions, so observed high quality spectra can be used as diagnostics of the geometry. Presently, it seems that active regions are favored over homogeneous slab coronae. (3) Observations of flux-index correlations can provide constraints on the source compactness. (4) The equivalent width of the iron line can be used as a diagnostic of the temperature of the reprocessed (black-body) radiation which then constrains the size." + }, + "9701/astro-ph9701112_arXiv.txt": { + "abstract": "We use Bayesian methods to study anisotropic models for the distribution of gamma ray burst intensities and directions reported in the {\\it Third BATSE Catalog} (3B catalog) of gamma ray bursts. We analyze data obtained using both the 64~ms and 1024~ms measuring timescales. We study both purely local models in which burst sources (``bursters'') are presumed to be distributed in extended halos about the Galaxy and M31, and mixed models consisting of a cosmological population of standard candle bursters and a local population distributed throughout a standard Bahcall-Soneira dark matter halo with a 2~kpc core. A companion paper studies isotropic models, including a variety of cosmological models, using the same methodology adopted here, allowing us to rigorously and quantitatively compare isotropic and anisotropic models. We find that the purely local models we have studied can account for the 3B data as successfully as cosmological models, provided one considers halos with core sizes significantly larger than those used to model the distribution of dark matter. A preference for cosmological over local models, or vice versa, must therefore be justified using information other than the distribution of burst directions and intensities. We infer core sizes for the halo distribution that are smaller than one might expectbased on popular semiquantitative arguments that consider the superposed dipole moments of shells centered on the Galactic center, and show why such arguments can lead to unwarranted conclusions. We also find that the 3B data do not constrain the width of power-law luminosity functions for burst sources. This disagrees with the findings of previous studies; we elucidate the qualitative reasons for the lack of a constraint, and discuss why our results differ from those of earlier studies. Our analysis of mixed models finds two families of models that can successfully account for the data: models with up to 20\\% of observed bursts in a bright local population visible to $\\sim 50$~kpc; and models with up to 50\\% of observed bursts in a dim local population visible only nearby (to less than a disk scale height). These models fit as well or better than purely cosmological models. They indicate that a surprisingly large local, anisotropic component could be present whose size is comparable to the sizes of hypothetical classes of bursts inferred from analyses of temporal and spectral characteristics. Finally, as in our study of isotropic models, we find substantial systematic differences between results based on 64~ms and 1024~ms data, indicating that a thorough understanding of the distribution of burst intensities and directions is likely to require detailed analysis of temporal properties. ", + "introduction": "At the time of the launch of the {\\it Compton Gamma Ray Observatory} ({\\it CGRO}), the prevailing viewpoint among gamma ray astrophysicists was that gamma ray bursts originate in the vicinity of neutron stars distributed throughout the Galactic disk (see, e.g., the reviews of Liang and Petrosian 1985, and Higdon and Lingenfelter 1990). Perhaps the strongest evidence supporting this hypothesis was the detection of absorption-like features at energies of $\\sim10$--50~keV in the spectra of bursts. The KONUS experiment provided the earliest evidence for the presence of such features (Mazets et al.\\ 1981, 1982), but the most conclusive evidence came from instruments on the {\\it Ginga} spacecraft (Murakami et al.\\ 1988; Fenimore et al.\\ 1988). The features observed by {\\it Ginga} have high statistical significance, and can be well-modelled as being due to cyclotron scattering in a strongly magnetized plasma with field strength $B\\sim 10^{12}$~G (Wang et al.\\ 1989; Lamb et al.\\ 1989). This field strength is typical of that associated with both rotation-driven and accretion-driven pulsars, suggesting that bursts are associated with strongly magnetized neutron stars. Further, requiring that a gravitationally confined scattering region be static implied that the burst sources so far observed were no further than a few hundred parsecs distant, otherwise the sources would have to be so luminous that radiation presure would drive the scattering plasma away from the source (Lamb, Wang, and Wasserman 1990; magnetic confinement may relax this constraint). The hypothesis that bursters formed a disk population seemed consistent with the most direct information available about the spatial distribution of bursters: the distribution of burst directions and intensities. The apparent isotropy of the distribution of directions to bursts (Atteia et al.\\ 1987; Golenetskii 1988; Hartmann and Epstein 1989; Hartmann and Blumenthal 1989) implied that we could be observing members of a disk population only to distances smaller than a disk scale height---the distance scale implied by magnetized neutron star models with a static cyclotron scattering region. Reconciling a disk population with the cumulative distribution of burst intensities (the ``size-frequency'' distribution) was somewhat more problematic. The cumulative distribution of burst fluences, $S$, or peak energy fluxes, $F$, was significantly flatter than the $-3/2$ power law expected from sampling a spatial distribution from well within its characteristic length scale, but it appeared that selection biases could account for much of the flattening (Yamagami and Nishimura 1986; Higdon and Lingenfelter 1986; Mazets and Golenetskii 1987; Paczynski and Long 1988; Schmidt, Higdon, and Hueter 1988). The burst peak count rate, $C$, was proposed as a less ``biased'' intensity measure than fluence or peak energy flux, and the distribution of peak count rates appeared to be consistent with a $-3/2$ power law. Dispute remained over whether the distribution began to flatten for the dimmest bursts (Jennings 1988), but the number of faint bursts was too small to ascertain whether such bursts were anisotropically distributed, as one would expect if the faint bursts were observed from beyond a disk scale height. Hartmann, Epstein, and Woosley (1990) modelled the distribution of neutron stars in the Galaxy, and found the direction and intensity observations to be consistent with an association of bursts with Population I neutron stars, provided the distribution was sampled to distances beyond $\\sim 150$~pc, but no greater than $\\sim 2$~kpc. Although the Galactic neutron star scenario appeared consistent with the observations, some investigators argued in favor of a cosmological origin for bursts (e.g., Usov and Chibisov 1975; van den Bergh 1983; Paczynski 1986; Goodman 1986). The combination of isotropy of burst directions and inhomogeneity implied by burst intensities is a natural characteristic of such models, provided the observations sample sources well beyond the local supercluster (Hartmann and Blumenthal 1989). An additional motivation for considering cosmological models was provided by Paczynski (1990) who, using a different model for the distribution of old neutron stars than that adopted by Hartmann, Epstein, and Woosley (1990), found a neutron star origin inconsistent with the distribution of burst directions and intensities. Together, the work of Paczynski and of Hartmann, Epstein, and Woosley implied that the consistency of the Galactic disk model with the observations depended on uncertain details of the models, particularly in regard to the distribution of birth velocities of pulsars and the detailed form of the Galactic potential (Hartmann, Epstein, and Woosley 1990). Six years prior to the launch of {\\it CGRO}, Meegan, Fishman, and Wilson (1985) reported detection of a single burst by a sensitive balloon-borne detector that should have seen $\\sim 43$ bursts if the population of burst sources was spatially uniform to the distance sampled by the detector, conclusively demonstrating that the cumulative distribution of intensities of dim bursts was flatter than the homogeneous $-3/2$ power law. This was thought to be consistent with the Galactic disk paradigm, provided that the detector was able to see bursts from sources more distant than a disk scale height, hence detecting inhomogeneity in the source distribution. It was thus predicted that the Burst and Transient Source Experiment (BATSE) on board {\\it CGRO} would find faint bursts concentrated in the Galactic plane, finally providing compelling evidence for the Galactic disk neutron star paradigm. Within a year of the launch of CGRO, BATSE observations spectacularly refuted the Galactic disk neutron star hypothesis (Meegan, et al.\\ 1992). The observations confirmed the inhomogeneity discovered with the earlier balloon observations: the cumulative distribution of the peak fluxes of BATSE bursts roughly follows a $-1$ power law and definitively rejects the $-3/2$ power law expected for a homogeneous distribution. Yet the distribution of the directions to these bursts is consistent with isotropy, with no significant concentration of burst sources in the Galactic plane. These basic features---apparent isotropy, and inhomogeneity---have been only more conclusively demonstrated by subsequent BATSE observations (Fishman et al.\\ 1992, 1996). Although the BATSE observations definitively rule out a Galactic disk origin for bursts, there is considerable controversy over whether the BATSE data favor cosmological models over local models that distribute burst sources in a large Galactic halo or corona, rather than in a disk population. Large scale isotropy and inhomogeneity are natural qualitative features of cosmological models for burst sources, so the BATSE observations have revitalized interest in such models. One can construct halo or coronal models that are consistent with the data, but they have length scales considerably larger than those for matter distributions known to be associated with the Galaxy prior to the BATSE observations, further fueling interest in cosmological models. But in the last five years, evidence has accumulated indicating that there may be a population of high velocity neutron stars with unbound or marginally bound orbits, very possibly forming a large Galactic corona (Lyne and Lorimer 1994; Frail 1996). Ironically, some of this evidence has been provided by observations of Soft Gamma Repeaters (SGRs; see, e.g., Rothschild 1996), another class of gamma ray transient studied with BATSE. This, combined with suggestive but so far unconclusive evidence for burst repetition (which is difficult to reconcile with most cosmological models) has revived interest in Galactic models, and enlivened the controversy over whether the data can discern between cosmological and local alternatives (see, for example, Lamb 1995 and Paczy\\'nski 1995). This paper is the third in a series in which we apply the principles of Bayesian inference to the problem of inferring the spatial and energy distribution of burst sources from burst direction and intensity data provided by BATSE. In Paper~I (Loredo and Wasserman 1995) we described the methodology and compared it to other methods in use. In Paper~II (Loredo and Wasserman 1996), a companion paper to this one, we apply the method to isotropic models (including cosmological models), using the data from the recently released {\\it Third BATSE Catalog} (Fishman et al.\\ 1996; hereafter the 3B catalog). In the present paper, we use the method to study anisotropic models. Our method is uniquely suited to the study of these models, because it is the only method presently available that is capable of analyzing the distribution of burst directions and intensities {\\it jointly}. Since all anisotropic physical models so far proposed have an anisotropy whose characteristics vary with burst intensity, only a method capable of analyzing the joint distribution can fully evaluate these models. In addition, since we are using the same method to study both isotropic and anisotropic models, we are able to quantitatively compare them. The Bayesian tool for doing this---the odds ratio---includes a factor that accounts for the size of the parameter spaces of models, resulting in an ``Ockham's Razor'' that automatically and objectively accounts for model complexity and parameter uncertainty in such comparisons. We presume the reader to be familiar with the notation and methodology described in Papers~I and II; \\S\\S~2 and 3 of Paper II summarize this information. The next section presents an analysis of halo models whose burst rate density is spherically symmetric about the Galactic center and falls off with radius $r$ like $1/[1 + (r/r_c)^2]$, where $r_c$ is a core radius parameter; we analyze models with and without a similar halo centered on M31. We study models with ``standard candle'' and power law luminosity functions. In \\S~3, we analyze models that superpose two populations: a standard candle halo population like those analyzed in \\S~2, and a standard candle cosmological population. For these models, we set the core size of the halo population equal to 2~kpc, the value inferred by Bahcall and Soneira (1980) in their study of the Galactic rotation curve. This allows us to obtain precise, model-dependent constraints on the fraction of bursts that could be associated with a known local matter distribution. We find this fraction can be quite high, even though the halo population has a relatively small core size. Finally, in \\S~4 we discuss some of the implications of the work reported here and in Paper~II. Before moving on to the results of our study, we note some additional distinguishing features of this study. No study has yet been published that has analyzed anisotropic models using the 3B data. This most recent BATSE catalog contains data for about twice as many bursts as the earlier 2B catalog (Fishman et al.\\ 1994), so new analyses using this data are of obvious importance. In addition, we analyze data for two of the three trigger timescales included in the catalog: 64~ms (the shortest) and 1024~ms (the longest). Previous studies used only one timescale; most studies used the 1024~ms timescale. As shown in Paper~II, the shapes of distributions of 64~ms and 1024~ms peak fluxes differ, and there is evidence that the additional structure present in the 1024~ms data is due to peak dilution (incorrect peak flux estimation when the peak duration is shorter than the measurement timescale). Thus it is important to analyze data from different timescales to ascertain what features of one's inferences are robust. Our method can be generalized to include temporal information about bursts, as outlined in Paper~I; but the required information is absent from the 3B catalog, and such an analysis is beyond the scope of the current investigation. Finally, we are not aware of a single study that calculated correct constraints for the unknown parameters of Galactic models, even using earlier BATSE data. Rather than using standard tools for calculating confidence regions, investigators instead computed goodness-of-fit statistics on grids throughout parameter space, and used contours of constant significance to constrain parameters (see, e.g., Hakkila et al.\\ 1994). This is not a correct parameter estimation technique, and the resulting ``significance regions'' are of no use beyond determining whether the best-fit model is acceptable. Since similar techniques have been used in other fields, we devote Appendix A to a general discussion of the problems with this approach. In particular, we apply it to a simple Gaussian estimation problem, and show that the methodology adopted by earlier investigators leads to grossly incorrect confidence regions, and that the error {\\it grows} with the size of the data set. By contrast, the Bayesian methodology produces probability densities for parameters of Galactic models for the 3B data, from which rigorous constraints on model parameters may be derived directly. ", + "conclusions": "Our analysis of halo models in \\S~2 demonstrates that such models can account for the 3B data as successfully as cosmological models, provided one considers halos with core sizes significantly larger than those used to model the distribution of dark matter. Only a few years ago, a local population of sources with such a large characteristic size would have been considered highly implausible; such a population would have to have been hypothesized purely for the purpose of hosting bursts. We now know that there is a population of high velocity neutron stars that could conceivably provide a host population with a very large characteristic length scale. Whether the characteristics of such a population could model the burst data as successfully as the {\\it ad hoc} halo models considered here is an open question, requiring detailed modelling beyond the scope of this study. Our work demonstrates how the analysis of such models can best be undertaken, and the results of \\S~2 should guide the study of other, more complicated local models. The core sizes we infer are smaller than one might expect based on popular semiquantitative arguments that consider the superposed dipole moments of shells centered on the Galactic center (Hartmann et al.\\ 1994; Briggs et al.\\ 1994). Such arguments are misleading, as we show in \\S~2. They fail to distinguish the artificial anisotropy arising from the displacement of the centers of Galactocentric shells from the observing point (the Sun), and the actual, intrinsic anisotropy best quantified by calculating angular moments of shells at a constant radius from the Sun. Our analysis of halo models also demonstrates that the 3B data do not constrain the width of power-law luminosity functions for burst sources. This result contradicts the findings of Hakkila et al.\\ (1995), who used less rigorous analysis methods and who restricted their search of parameter space to a significantly smaller region than that explored here. We elucidate the qualitative reasons for the lack of such a constraint in \\S~2. As with the isotropic models studied in Paper~II, inferences based on the 64~ms and 1024~ms data are formally inconsistent, in the sense that there is negligible overlap of the credible regions found by analyzing the two data sets. Although the shapes of the burst distributions that best model each data set differ somewhat, the inconsistency arises largely because the two data sets imply very different burst rates per unit volume. The 64~ms data implies rates about 40\\% larger than the 1024~ms data, and in this sense the 64~ms timescale is {\\it more} sensitive than the 1024~ms timescale, even though the 1024~ms data set is larger (i.e., the 1024~ms data set is not as large as one would expect from extrapolating the 64~ms data set to the lower fluxes detectable using the 1024~ms timescale). In Paper~II we argue that this difference could arise from ``peak dilution'' in the 1024~ms data set: when bursts have peaks briefer than 1024~ms, their intensities are underestimated by 1024~ms measurements. This changes the shape and normalization of the intensity distribution in a manner that may account for the discrepancy between the two data sets. The 64~ms data are not immune to such effects, but will be less affected. We also studied two-population models, consisting of superposed standard candle cosmological and local halo populations. Numerous studies suggest that there may be two (or more) classes of bursts, as we review in \\S~3. These models also serve a purely pragmatic purpose of allowing precise, model-dependent quantification of the constraints the 3B data place on the anisotropy of the distribution of burst sources. They also provide a model-dependent test of the adequacy of cosmological models. For the two-population models, we took the halo population to follow the distribution of dark matter in a standard Bahcall-Soneira halo with a core size of 2~kpc, so that the local population is associated with a known distribution of matter. Two families of models successfully account for the data: models with luminous halo sources visible to $\\sim 50$~kpc; and models with dim halo sources visible from within a disk scale height. Despite the fact that the luminous halo sources would be distributed anisotropically, models with $\\approx 10$\\% of observable bursts from the halo are favored, and halo fractions as large as 20\\% are acceptable. Dim halo sources would comprise a nearly isotropic, homogeneous component. For such sources, the data favor large halo fractions, of the order of 40\\% to 50\\%. These results are consistent with the relative sizes of classes of bursts inferred from characteristics of burst lightcurves, or with the fraction of bursts observed by {\\it Ginga} to have low energy absorption features. We have not yet ascertained whether membership in these classes is correlated with burst intensity in the manner that these models would predict. The common methodology employed here and in Paper~II, where we analyze cosmological models, permits us to rigorously compare how well cosmological and local models account for the full joint distribution of burst peak fluxes and directions. The Bayesian tool for this comparison---the Bayes factor---objectively accounts for parameter uncertainty, providing a quantitative ``Ockham's Razor.'' We find that the data do not decisively prefer any one of the models we have studied to its competitors. In particular, local models and models with a substantial local component account for the 3B burst intensity and direction data as well as purely cosmological models. Despite this, several investigators strongly prefer cosmological models to local ones, or vice versa. This preference can be justified only by consideration of information beyond that in the distribution of burst peak fluxes and directions analyzed here. Numerous studies have been undertaken of some additional burst characteristics, including searches for evidence of time dilation, burst repetition, and spectral lines. To date, such searches have been inconclusive, with investigators who use different frequentist methodologies arriving at different, often conflicting conclusions. The Bayesian approach we have adopted in this work can be straightforwardly generalized to address many of these issues. Some such generalizations have already been outlined in Paper~I, and we are pursuing studies of some of these outstanding controversial issues from within this framework." + }, + "9701/astro-ph9701177_arXiv.txt": { + "abstract": "We propose a method to determine the cosmic mass density $\\Omega$ from redshift-space distortions induced by large-scale flows in the presence of nonlinear clustering. Nonlinear structures in redshift space such as fingers of God can contaminate distortions from linear flows on scales as large as several times the small-scale pairwise velocity dispersion $\\sigv$. Following Peacock \\& Dodds (1994), we work in the Fourier domain and propose a model to describe the anisotropy in the redshift-space power spectrum; tests with high-resolution numerical data demonstrate that the model is robust for both mass and biased galaxy halos on translinear scales and above. On the basis of this model, we propose an estimator of the linear growth parameter $\\beta = \\Omega^{0.6}/b$, where $b$ measures bias, derived from sampling functions which are tuned to eliminate distortions from nonlinear clustering. The measure is tested on the numerical data and found to recover the true value of $\\beta$ to within $\\sim 10$\\%. An analysis of \\iras\\ galaxies yields $\\beta = 0.8^{+0.4}_{-0.3}$ at a scale of 1,000~\\kms\\ which is close to optimal given the shot noise and the finite survey volume. This measurement is consistent with dynamical estimates of $\\beta$ derived from both real-space and redshift-space information. The importance of the method presented here is that nonlinear clustering effects are removed to enable linear correlation anisotropy measurements on scales approaching the translinear regime. We discuss implications for analyses of forthcoming optical redshift surveys in which the dispersion is more than a factor of two greater than in the IRAS data. ", + "introduction": "A redshift-space map of galaxies is distorted relative to the real-space galaxy distribution as a result of peculiar motions along an observer's line of sight. These distortions generate an anisotropy in pairwise correlations which would not appear in the real-space galaxy distribution of a statistically homogeneous and isotropic universe. Kaiser (1987) quantified the correlation anisotropy that results from large-scale peculiar flows in terms of the power spectrum of galaxies using the linear theory of gravitational instability. He demonstrated that in the linear regime the power measured by a distant observer only depends on the angle between the wavevector and the observer's line of sight, and on the dimensionless factor $\\beta \\equiv \\Omega^{0.6}/b$. Here, $\\Omega$ is the cosmic mass density parameter and $b$ is a function that differs from unity if the galaxies are a biased sample of the total mass. Hamilton (1992) transformed Kaiser's result out of the Fourier domain to determine the redshift-space correlation function $\\xis$ in linear theory and proposed an estimator of $\\beta$ based on a spherical harmonic decomposition of $\\xis$. Subsequently, a number of $\\beta$ measurements from linear flow distortions have been reported (e.g., Hamilton 1993, 1995; Bromley 1994; Fisher, Scharf \\& Lahav 1993; Fisher et al. 1994; and Nusser \\& Davis 1994). There are two major challenges for a statistical measurement of $\\beta$ from linear flows in a real galaxy redshift catalog. The first comes from the finite size of the catalog. The catalog must contain a fair sample of structure on a particular scale or else the correlation information will contain finite-sampling noise. A reasonable criterion is that sampling scales should be below $\\sim 10$\\% of the characteristic size of survey volume. The second challenge is to ensure that the sampling scales are large enough so that signal from nonlinear clustering does not contaminate the linear fluctuation modes. Fingers of God can extend up to several thousand kilometers per second in optically selected galaxy surveys and clearly affect the redshift-space power on these scales. (Here redshift-space distances are usually given in terms \\kms; where real-space lengths are needed, we use the conventional units of \\hmpc\\ where $h$ is the Hubble parameter in units of 100~\\kms.) Thus the signal in redshift space from linearly growing fluctuation modes on scales from $\\sim$1,000~\\kms\\ (characteristic of the translinear regime below which linear theory breaks down) to $\\sim$4,000~\\kms\\ can be contaminated by strongly nonlinear features. Present-day catalogs are typically within a few times 10,000~\\kms\\ in radial extent, leaving at best only a small range of scales on which purely linear modes can be measured with a high degree of statistical integrity. In this paper we address this problem of determining $\\beta$ from linear flows in the presence of nonlinear clustering. Our approach is motivated by a remarkable result from Peacock \\& Dodds (1994) concerning the redshift-space power spectrum $\\ps$. While $\\xis$ itself can be modeled given the distribution of pairwise velocities, there is evidence (Fisher et al. 1994; Warren 1995) that this distribution has complicated behavior on a critical range scales from $\\sim$1--10~\\hmpc. In contrast, Peacock \\& Dodds (1994) found that nonlinear signal can be modeled simply and accurately in the Fourier domain on translinear scales and above. Specifically, the effects nonlinear clustering and linear flows in redshift space can be described by separable linear filters acting on the real-space power spectrum. The linear filtering hypothesis has been supported by tests with $N$-body simulations, although only mass particles and not collapsed galaxy halos were considered (Gramann, Cen \\& Bahcall 1993; Tadros \\& Efstathiou 1996). Here, we first suggest a form of the filters, which amounts to constructing a model for the anisotropy in $\\ps$, and determine its validity in high-resolution numerical simulations using both mass particles and collapsed halos which are identified as galaxies. On the basis of this model we then suggest a method to extract $\\beta$ on scales of 1,000~\\kms\\ and above, even when the velocity dispersion is of comparable size. Finally, we apply our method to the \\iras\\ survey. The result is a new measure of $\\beta$ on translinear scales. ", + "conclusions": "Here we have modeled the anisotropy in the redshift-space power spectrum on the basis of Peacock \\& Dodds' (1994) theoretical work and our own high-resolution numerical simulations. We demonstrated that the effects of nonlinear clustering on the power spectrum of both mass and galaxy halos can be described by a simple multiplicative filter function. The filter-based model for the anisotropy is seen to hold on linear and translinear scales which are greater than the pairwise velocity dispersion $\\sigv$ at megaparsec. Having established the domain of validity, we then proposed a redshift-space statistic, the ratio of sample variances in equation~(\\ref{eq:omrat}), which is both impervious to the effects of nonlinear clustering and sensitive to linear flows and the linear growth parameter $\\beta$. This work should serve as a cautionary reminder of the extent to which the nonlinear distortions contaminate signal from linear flows (e.g., Fig.~2). In a measurement of power, the magnitude of contamination can be determined from equation~(\\ref{eq:pfilt_disp}); for example the nonlinear signal reduces the power along the line of sight by $\\sim 40$\\% at scales $2\\pi/k$ equal to five times $\\sigv$. In the case of the \\iras\\ survey where $\\sigv$ is $\\sim 300$~\\kms, the presence of nonlinear distortion is important even at scales of 1,500~\\kms. For optical surveys where $\\sigv \\sim 800$~\\kms\\ (e.g., Marzke et al. 1995), the extent of contamination exceeds $4,000$~\\kms. The work presented here also gives the encouraging message that nonlinear clustering does not prevent the accurate extraction of $\\beta$ from linear flows in redshift data. Figure~2 provides the evidence based upon cosmological simulations of the greatest dynamical range published to date. The analysis of the \\iras\\ survey is intended to show that the method introduced here gives reasonable results. Our estimate of $\\beta = 0.8^{+0.4}_{-0.3}$ is consistent at the 2-$\\sigma$ level with virtually every published value based on an analysis of the IRAS galaxies (e.g., Dekel et al. 1993; Fisher, Scharf \\& Lahav 1994; Cole, Fisher \\& Weinberg 1995; Hamilton 1995), although it is higher than most redshift-space measurements based exclusively on linear theory. Our procedure is most similar in spirit to that of Hamilton (1995) who obtained $\\beta = 0.69^{+0.21}_{-0.19}$ using a merged \\iras\\ and the optically selected QDOT catalog. Hamilton, working with harmonics of a smoothed power spectrum, modeled the nonlinear noise in the same way as we did, with an isotropic exponential distribution for pairwise velocities. We hope that we have provided convincing evidence in \\S2 that such a model is indeed robust. We note that our estimate of $\\beta$ from the \\iras\\ sample is consistent with dynamical analyses that make use of real-space information as well as redshift data. For example, Dekel et al. (1993) find $\\beta = 1.28 \\pm 0.3$ on scales of 1,500--4,000~\\kms. While knowledge of both real-space and redshift distributions can yield a more accurate measurement of $\\beta$, the real-space data are difficult to measure and therein lies the value in extracting $\\beta$ directly from redshift space. Fisher et al. (1994) performed a pure redshift-space analysis of the \\iras\\ galaxies by modeling $\\xis$ and obtained $\\beta = 0.45^{+0.27}_{-0.18}$ on scales of 1,000--1,500~\\kms, well below the dynamical estimates at larger scales. However, the modeling of $\\xis$ and its dependence on the uncertain behavior of the velocity distribution function is difficult. In contrast, our measure of $\\beta$, based instead on a much simpler model in the Fourier domain, is consistent with the large-scale dynamical results. The ultimate hope is that our measure may be applied to forthcoming optical surveys which exhibit high $\\sigv$ values (e.g., Marzke et al. 1995). With the ability to accurately determine $\\beta$ down to translinear scales in spite of nonlinear clustering effects, our measure can provide high statistical significance in a correlation analysis of redshift-space data." + }, + "9701/astro-ph9701031_arXiv.txt": { + "abstract": "s{Cosmic microwave anisotropy satellites promise extremely accurate measures of the amplitude of perturbations in the universe. We use a numerical code to test the accuracy of existing approximate expressions for the amplitude of perturbations produced by single-field inflation models. We find that the second-order Stewart--Lyth calculation gives extremely accurate results, typically better than one percent. We use our code to carry out an expansion about the general power-law inflation solution, providing a fitting function giving results of even higher accuracy.} ", + "introduction": " ", + "conclusions": "" + }, + "9701/astro-ph9701207_arXiv.txt": { + "abstract": "We have developed an excitation model for small molecules including radiative pumping by dust photons from near infrared to submillimeter wavelengths. This model applies to molecules within bright photodissociation regions in galactic star-forming regions and starburst galaxies. In such environments, the near infrared photons emitted by polycyclic aromatic hydrocarbon molecules (PAHs) or small carbon grains, $5 \\mic < \\lambda < 20 \\mic$, are able to penetrate molecular regions and pump the molecules in an excited vibration state. The far infrared and submillimeter transitions involved in the de-excitation cascade are strongly affected by this process. Their intensities can be enhanced by several orders of magnitude. We have applied this model to $\\HtwoO$ and $\\NHtrois$, ortho and para species. The behaviour of the FIR rotation lines with respect to the gas density and the molecule column density is strongly modified compared to pure collisional models. At moderate densities, $\\nHdeux < 10^5 \\cmcube$, radiative pumping dominates the excitation of the rotation ladder, and the FIR lines are good probes of the molecule column density. The near infrared absorption and emission lines predicted to appear between 6 and 7 $\\mic$ for $\\HtwoO$ and from 10 to 12 $\\mic$ for $\\NHtrois$ can also be used for this purpose. Concerning $\\orthoHtwoO$, a quasi resonant pumping of the $\\nu_2$ $6.18 \\mic$ transition by photons of the 6.2 $\\mic$ PAH band occurs. A strong de-excitation emission line is expected at 6.64 $\\mic$. ", + "introduction": "Modelling the formation of molecular rotation lines in interstellar clouds is a crucial task both for theoretical and observational astrophysics. On the theoretical side, early studies by Goldreich and Kwan (\\cite{Goldreich74}) and Goldsmith and Langer (\\cite{Goldsmith78}) have shown that the rotation lines from polar and/or abundant molecules like $CO$, $\\Otwo$, $HD$, $\\HtwoO$ and other hydrides should be the major cooling agents of the molecular gas. Together with the heating mechanisms (e.g. cosmis rays, $\\Htwo$ formation, gravitational contraction and grain heating), they control the thermal balance of dense clouds, and thus, the conditions that will allow the contraction toward star formation. On the observational side, millimeter and submillimeter molecular lines are used to probe several physical parameters of molecular clouds, such as the gas temperature, the $\\Htwo$ density, the abundance of molecules, and the gas dynamics. Densities and temperatures in molecular clouds are such that the collision rate with $\\Htwo$ molecules generally does not allow to reach the thermodynamic equilibrium of the rotation ladder (see Goldreich and Kwan \\cite{Goldreich74}). Instead, the coupling with the radiation field is important, both via absorption and spontaneous emission. Takahashi et al. (\\cite{Takahashi83} and \\cite{Takahashi85}), have shown how far infrared dust photons, $\\lambda = 30$ to $100 \\mic$ are able to pump the rotation ladder of the $\\HtwoO$ molecule within warm molecular cores. Carroll and Goldsmith (\\cite{Carroll81}) have investigated the pumping of molecules via their near infrared (NIR) vibration transitions. At that time it was assumed that a significant fraction of the dust emission could be in the NIR range, only in the immediate vicinity of hot stars. However, the effect of the NIR radiation attributed to large polycyclic aromatic hydrocarbon molecules (PAHs) and very small grains, which makes up to 30 \\% of the total dust emission (Puget et al. \\cite{Puget85}), has never been taken into account. This radiation is characterised by a strong continuum and broad emission features, located at 6.2 and 7.7 $\\mic$ for the strongest ones, wavelengths which co\\\"{\\i}ncide with the vibration transitions of some molecules. For instance the $v2$ transition of $\\HtwoO$ and the $v4$ transition of $\\NHtrois$ occur at the maximum of the 6.2 $\\mic$ PAH feature, and the 7.8 $\\mic$ vibration of CS falls within the wavelength range of the 7.7 $\\mic$ feature. There is now clear observational evidence that the NIR emission bands are not only present in peculiar reflexion nebulae, but are a general component of the dust radiation: detection of the 3.3 and 6.2 $\\mic$ bands in our galaxy by Giard et al. \\cite{Giard94a} and Ristorcelli et al. \\cite{Ristorcelli94}, and, more recently, the measurements performed by the ISO satellite (Boulanger et al. \\cite{Boulanger96}, Mattila et al. \\cite{Mattila96}, and Verstraete et al. \\cite{Verstraete96}). The aim of this paper is to revisit the problem of the coupling between molecules and the dust radiation field, in the framework of a model which includes the PAH emission bands. The mechanism, and the rates involved in some representative astrophysical environments, are presented in Sect. \\ref{NIRpump}. A detailed model for the molecular excitation in star-forming photodissociation regions (PDR) is develloped in Sect. \\ref{Model}. The results for two molecules, $\\HtwoO$ and $\\NHtrois$, are presented in Sect. \\ref{Results}. ", + "conclusions": "Conclusion} The main conclusions that can be drawn from our study are the following: 1) NIR photons are able to contribute very significantly to the rotational excitation of polar molecules within PDRs in galactic star-forming regions and starburst galaxies. 2) As this mechanism is independent of the gas density, we predict that bright FIR and submillimeter rotation lines should be observed from the relatively moderate density and non shocked gas in PDRs, $\\nHtwo < 10^5 \\cmcube$, where molecules are expected to be abundant because of grain mantles evaporation. 3) For $\\HtwoO$ and $\\NHtrois$ the intensity of the FIR and submillimeter rotation lines can be increased by a factor ranging from 2 to several 100 with respect to models which do not take into account NIR pumping. This means that the molecular column densities that can be infered from observations toward such regions, using standard models, can be wrong by similar factors. 4) For gas densities below $10^5 \\cmcube$, where NIR pumping dominates over collisional excitation, the intensities of the FIR lines becomes mostly independant of the gas density. This implies that the FIR and submillimeter lines can be used to probe the molecule column density if the radiation field has been correctly estimated. 5) For $\\orthoHtwoO$ we have a quasi-resonant excitation of the molecule, by absorption of 6.2 $\\mic$ PAH photons in the 6.18 $\\mic$ fundamental vibration transition. We thus predict a high line to continuum contrast for the emission de-excitation line at 6.64 $\\mic$.\\\\ NIR pumping is likely to affect not only the PDRs, as modelled in this paper, but also the molecular clouds and cores in their vicinity. The extinction in the 6 to 8 $\\mic$ range is minimum (typically 0.02 of visible extinction, Rieke and Lebofsky \\cite{Rieke85}) and thus, the radiation produced in the PDR penetrates well in the molecular clouds. Most molecules present in PDRs and nearby molecular clouds will actually be affected by radiative NIR pumping. In addition to $\\HtwoO$ and $\\NHtrois$ one can list $CS$, $HCN$, $CO_2$, $CH_4$, $H_3O^+$, $HC_3N$, and $CH_3OH$. The combination of the absorption and emission lines from all species existing in the gas phase is likely to be responsible for the very complex detailed structure of the 5 to 20 $\\mic$ spectra measured at moderate spectral resolution in the direction of the M17-SW PDR with the ISO-SWS intrument (Verstraete et al. \\cite{Verstraete96}). Decrypting such spectra in terms of molecular abundances will require the production of synthetic model spectra including a significant number of molecules of astrophysical interest. This will be done in a forthcoming paper." + }, + "9701/astro-ph9701082_arXiv.txt": { + "abstract": "We assess the prospects for attaining steady nuclear flow equilibrium in expanding $r$-process environments where beta decay and/or neutrino capture determine the nuclear charge-changing rates. For very rapid expansions, we find that weak steady flow equilibrium normally cannot be attained. However, even when neutron capture processes freeze out in such nonequilibrium conditions, abundance ratios of nuclear species in the $r$-process peaks might still {\\it mimic} those attained in weak steady flow. This result suggests that the $r$-process yield in a regime of rapid expansion can be calculated reliably only when all neutron capture, photodisintegration, and weak interaction processes are fully coupled in a dynamical calculation. We discuss the implications of these results for models of the $r$-process sited in rapidly expanding neutrino-heated ejecta. ", + "introduction": "In this paper we study the influence of weak charge-changing nuclear reactions (beta decay and neutrino capture) on neutron capture nucleosynthesis in rapidly expanding (decompressing) media. The astrophysical site of origin of $r$-process nucleosynthesis (or {\\it rapid neutron capture} nucleosynthesis; Burbidge {\\it et al.} 1957; Cameron 1957) is not known with certainty ({\\it cf.} Mathews \\& Cowan 1990). Recently, however, there has been a resurgence of interest in $r$-process nucleosynthesis. This interest stems from new observations of heavy elements in low metallicity halo stars (Sneden {\\it et al.} 1996) and the possibility that the $r$-process may be sited in neutrino-heated supernova ejecta (Meyer {\\it et al.} 1992; Woosley \\& Hoffman 1992; Takahashi, Witti, \\& Janka 1994; Woosley {\\it et al.} 1994). These considerations could have implications for particle physics and cosmology since, if we understand the abundance yields of decompressing neutrino-heated material, then it is possible that we could probe the basic properties of neutrinos (Fuller {\\it et al.} 1992; Qian {\\it et al.} 1993; Qian \\& Fuller 1995). Though models of $r$-process nucleosynthesis from neutrino-heated ejecta are promising, they suffer from a number of flaws. For example, it is not understood how the neutron-to-seed nucleus ratio in neutrino-heated supernova ejecta can become high enough to produce an abundance pattern yield which will match that of the solar system (see {\\it e.g.} Hoffman, Woosley, \\& Qian 1996; Meyer, Brown, \\& Luo 1996). Yet it may be required that at least {\\it some} supernovae produce a solar $r$-process distribution, as there may be direct observational evidence that even the earliest $r$-process events in the galaxy produced an abundance pattern consistent with that observed in the solar system (see {\\it e.g.} Sneden {\\it et al.} 1996). An alternative $r$-process site that may not suffer from the neutron-to-seed nucleus problem is the decompression of \\lq\\lq cold\\rq\\rq\\ neutron matter from neutron star collisions (Lattimer {\\it et al.} 1977; Meyer 1989). Despite problems with these models, it may be {\\it necessary} to consider $r$-process environments sited in an intense neutrino flux. It has been argued that the observed solar system $r$-process abundance pattern itself may contain clues that point to $r$-process nucleosynthesis occuring in an intense neutrino flux (McLaughlin \\& Fuller 1996) or requiring a significant neutrino fluence (Qian {\\it et al.} 1996; Haxton et al. 1996). Neutrino post-processing effects were also discussed in Meyer et al. (1992). The neutrino-heated supernova ejecta environment and the neutron star merger site would each suggest that rapid neutron capture nucleosynthesis takes place in an expanding medium and in an intense neutrino flux. In this paper we concentrate on how the nuclear flows in a rapid neutron capture environment can be influenced by the interplay of weak charge changing reactions and rapid material outflow. Our study is meant to extend the evaluations of the \\lq\\lq waiting point\\rq\\rq\\ assumption (that neutron captures are balanced against photodisintegrations along an isotopic (constant $Z$) chain, see {\\it e.g.}, Cameron, Cowan, \\& Truran 1982) to the unique conditions of rapid outflow and intense neutrino flux which may characterize neutrino-heated supernova ejecta or decompressing neutron star matter. The intense neutrino flux provides a new wrinkle for $r$-process calculations: in addition to the usual beta decay processes, electron neutrino $\\nu_e$ captures on nuclei can change nuclear charge (Nadyozhin \\& Panov 1993; Fuller \\& Meyer 1995; McLaughlin \\& Fuller 1995, 1996; Qian 1996; Qian {\\it et al.} 1996). Further complicating matters, neutrino capture rates can be position (time) dependent, unlike beta decay rates. In the comoving frame of a fluid element rapidly receding from a neutrino source, the neutrino capture rate on heavy nuclei will fall sharply with time. Indeed, it has been suggested that rapid outflow may allow neutrino capture to dominate over beta decay at the onset of the neutron capture epoch, yet allow beta decays to dominate weak flows after neutron capture ceases in the regime where neutron-rich nuclei \\lq\\lq decay back\\rq\\rq\\ toward the valley of beta stability (McLaughlin \\& Fuller 1996; Qian 1996; Qian {\\it et al.} 1996). Weak nuclear flows help determine the final abundance ratios of nuclear species in the abundance peaks. It has been argued that the solar system data provides evidence that these ratios are within $20 \\%$ of the predictions of calculations based on weak steady flow (Kratz {\\it et al.} 1988). In weak steady flow, all abundances remain constant in time: neutron capture/photodisintegration or $(n,\\gamma )-(\\gamma , n)$ equilibrium obtains for the species along an isotopic $Z$ chain, and weak flows couple the abundances of adjacent isotopic chains. We can identify three regimes which characterize weak flows in r-process nucleosynthesis in a rapidly expanding medium. These are: (1) the case where the expansion timescale (or, more precisely the time available for neutron capture) for the medium is long compared to the weak charge-changing timescale (given by the inverse of the sum of the typical neutrino capture and beta decay rates); (2) the case where the expansion timescale is short compared to the weak charge-changing timescale and (3) the case where these timescales are comparable. We will show that weak steady flow equilibrium can only be guaranteed in case (1), while it is impossible in case (2) and doubtful in case (3). In Section 2 we consider these cases and discuss the prospects for attainment of weak steady flow. We explicitly integrate the differential equations for the abundances of the nuclear species in the neutron number $N=82$ peak under several restrictive assumptions, including that of instantaneous $(n,\\gamma )-(\\gamma , n)$ equilibrium. In Section 3 we employ these calculations to produce a plot suggesting which regions of expansion timescale and neutrino flux parameter space might be conducive to either attaining weak steady flow equilibrium or abundance ratios which mimic those derived from weak steady flow. In Section 4 we assess the implications of these results for models of $r$-process nucleosynthesis from neutrino-heated ejecta in general and wind models of the post-core-bounce supernova environment in particular. ", + "conclusions": "In this paper we have presented a guide for analyzing the weak charged changing flow in a potential $r$-Process environment. We have identified three different regimes for the weak flow, which depend on the neutrino flux and the expansion timescale. Steady weak flow occurs when the time spent in $(n,\\gamma)-(\\gamma,n)$ equilibrium is very long when compared with the inverse charge changing rate. In order to asses the impact of these weak flow considerations on the nucleosynthesis yield in a potential $r$-Process site, in section 3 we applied the weak flow analysis to a model with an exponential outflow. We concentrated on the $N=82$ peak, although the same analysis may also be applied to other regions such as the $N=126$ peak. (Clearly, the favored parameter space of expansion timescale and neutrino flux will be somewhat different for the material which at the conclusion of neutron capture makes the $A = 195$ material.) We find that the material is only in weak steady flow equilibrium when the expansion timescale is very long, such as a few seconds. At the much shorter times favored in the wind model, such as a tenth of a second, the $r$-process abundances can still conceivably mimic those of steady beta flow. However, these probably finely tuned scenarios might necessitate neutrino-induced neutron spallation post-processing during the decay back to beta stability at a level which could be intolerable. A full network calculation carried through the period of decay back to beta stability would be necessary, however, to answer this question definitively. {}From our analysis of steady weak flow, we can infer the sort of conditions which would be conducive to producing the peaks in the $r$-process abundance distribution. One possibility is that neutrino capture has a limited impact during the neutron capture phase. Then the potential environments are restricted to those where steady beta flow obtains or nearly obtains. One option might be a slightly altered version (altered to produce the correct neutron-to-seed ratio) of the slow one dimensional post-core-bounce outflow produced in the Type II supernova model of Mayle and Wilson (as employed in the Meyer {\\it et al.} 1992 and Woosley {\\it et al.} 1994 $r$-process calculations). Another possibility is that the neutrino captures play a major role in accelerating the weak charge changing flow during the neutron capture phase, but the expansion rate takes on a more complicated time dependence than the single timescale exponential wind model employed here. An example of such a scenario occurs when the expansion rate is fast during alpha rich freeze out so that the neutron-to-seed ratio is acceptable ({\\it e.g.} Hoffman, Woosley, \\& Qian 1996), slow during the period of $(n,\\gamma)-(\\gamma,n)$-equilibrium (neutron capture epoch ) so that steady weak flow can obtain, and then relatively fast during the post-processing phase so that the amount of neutron spallation is restricted. Such a \\lq\\lq fast-slow-fast\\rq\\rq\\ scenario sounds excessively convoluted in the one-dimensional post-core-bounce outflow models, but it remains to be seen whether three dimensional convective models could produce such conditions (envision initially rapidly outflowing material which during the epoch of neutron capture is caught in a convective eddy and then later re-accelerated and ejected). A third possibility is that the site of $r$-Process nucleosynthesis is not the post-core-bounce supernova environment. We caution, however, that decompression of cold neutron matter from neutron star mergers/collisions could well experience the same parameter space \\lq\\lq squeeze\\rq\\rq\\ we describe here for wind models of neutrino-heated supernova ejecta. We conclude by emphasizing that only a full network calculation, including the neutron capture period and the post processing period can predict accurately the nucleosynthesis yield which results from a rapidly expanding environment." + }, + "9701/astro-ph9701119_arXiv.txt": { + "abstract": " ", + "introduction": "Photometry in two distant clusters of galaxies by Butcher and Oemler (1978) revealed a surprisingly large population of blue galaxies (Butcher-Oemler effect). Follow-up spectroscopy by Dressler and Gunn (1982, 1983 [DG83]) and Dressler, Gunn, and Schneider (1985 [DGS]) were pioneering steps towards confirming cluster membership and unraveling their nature from spectral lines. Among six blue galaxies observed in 3C 295, three appeared to be active galactic nuclei (AGN) and three to have strong Balmer absorption lines and negligible emission lines suggestive of a burst of star formation (DG83). Among 14 blue members observed in Cl 0024+16, three also appeared to be AGNs, but the remaining resembled local spirals with extended star formation, rather than starbursts as in 3C295 (DGS). Since their line-ratio criterion for AGNs was not foolproof and their spectra were of low spectral resolution (over 1000 km-s$^{-1}$), DGS suggested that linewidths would be a useful check of their claimed AGNs. This {\\it Letter} presents such linewidths as well as new line ratios and results of searches for blue compact nuclei for all three AGN candidates in Cl 0024+1654 as well as for three other emission line galaxies. None are confirmed to be AGNs. We adopt $h = 0.5$, $q_0 = 0.05$, and $\\Lambda = 0$, i.e., $\\Omega_0 = 0.1$, as our cosmology. Given these parameters, $L^*$ ($M_B \\sim -21$) corresponds to $r \\sim 21.0$ and $1\\arcsec$ spans 7.0 kpc at the cluster redshift of $z = 0.4$. ", + "conclusions": "" + }, + "9701/astro-ph9701096_arXiv.txt": { + "abstract": "Variation in the angle $\\alpha$ between a pulsar's rotational and magnetic axes would change the torque and spin-down rate. We show that {\\em sudden} increases in $\\alpha$, coincident with glitches, could be responsible for the persistent increases in spin-down rate that follow glitches in the Crab pulsar. Moreover, changes in $\\alpha$ at a rate similar to that inferred for the Crab pulsar account naturally for the very low braking index of the Vela pulsar. If $\\alpha$ increases with time, all pulsar ages obtained from the conventional braking model are underestimates. Decoupling of the neutron star liquid interior from the external torque cannot account for Vela's low braking index. Variations in the Crab's pulse profile due to changes in $\\alpha$ might be measurable. ", + "introduction": "Although the rotational behavior of pulsars has been monitored in detail for decades, some aspects of how an isolated pulsar spins down are unresolved. Questions surround the manner in which angular momentum is removed from the system and how the different components of the neutron star interior couple to one another. In particular, most pulsars do not slow in a regular fashion, but undergo variations in their spin rates in the form of glitches and timing noise. Except for such timing irregularities, a pulsar is expected to slow down steadily. For example the {\\em vacuum dipole model}, a theory of rotational energy loss to magnetic dipole radiation (\\cite{dipole_model}), predicts $\\dot{\\Omega}\\propto -\\Omega^3$, where $\\Omega$ is the pulsar rotational velocity. This spin-down law gives a braking index of $n_{\\rm obs}\\equiv \\Omega\\ddot{\\Omega}/\\dot{\\Omega}^2=3$. Timing irregularities make meaningful determination of $\\ddot{\\Omega}$ and hence $n_{\\rm obs}$ difficult or impossible in most cases. Consequently, the only braking indices available until recently were those for three very young, relatively quiet pulsars: the Crab pulsar ($n_{\\rm obs}=2.51\\pm 0.01$; \\cite{LPS}), PSR B1509-58 ($n_{\\rm obs}=2.837\\pm 0.001$; \\cite{psr1509_index}), and PSR B0540-69 ($n_{\\rm obs}=2.24\\pm 0.04$; \\cite{psr0540_index}). These braking indices are in significant disagreement with the vacuum dipole model expectation. Plasma in the pulsar magnetosphere could carry angular momentum from the star and alter the magnetic structure from the dipolar configuration, giving a braking index as low as one (see, \\eg, \\cite{wind}). Another possibility is that the magnetic moment of the star changes in time, through either a change in the surface field strength (see, \\eg, \\cite{field_growth}; \\cite{MP}) or the angle between the magnetic and spin axes (see, \\eg, \\cite{counteralignment}; \\cite{LEB}; \\cite{PP}). In the Crab pulsar, persistent increases in the spin-down rate are observed to accompany glitches (see, \\eg, \\cite{LPS} and Table 1). This phenomenon can in principle be explained by either a sudden change in the external torque or in the moment of inertia acted upon by the torque. In this Letter we suggest that the persistent increases in the Crab's spin-down rate reflect sudden, glitch-induced reorientations of the star's magnetic axis that increase the external torque. Recently, Lyne \\etal\\ (1996) attempted to separate the glitch activity of the Vela pulsar from its underlying spindown, and obtained a surprisingly small value of $\\ddot{\\Omega}$. If this small $\\ddot{\\Omega}$ reflects the true spin-down of the star, the braking index is only $1.4\\pm 0.2$. This braking index, far smaller than those for younger pulsars, suggests that the braking index decreases with time. This possibility forces reconsideration of traditional determinations of pulsar ages that assume constancy of the pulsar magnetic moment. In this Letter we demonstrate that shifts in $\\alpha$ at a rate similar to that inferred for the Crab account naturally for Vela's low braking index. ", + "conclusions": "We have shown that the persistent offsets in the Crab's spin-down rate following glitches could be due to sudden glitch-induced increases in the angle $\\alpha$ between the rotational and magnetic axes. Moreover, a similar average growth rate of $\\alpha$ accounts for the very low braking index of the Vela pulsar. In the vacuum dipole model, the characteristic growth times of $\\alpha$ in these two pulsars are $\\tau_\\alpha\\sim 7\\times 10^4$ yr for the Crab and $\\tau_\\alpha\\sim 3\\times 10^4$ yr for Vela. Lyne and Manchester (1988) found that pulsars older than $\\sim 10^6$ yr tend to have smaller $\\alpha$ than younger pulsars. If our suggestion that $\\alpha$ grows in young pulsars is correct, and the trend found by Lyne and Manchester (1988) is statistically significant, then the evolution of $\\alpha$ is not monotonic. That is, in young pulsars ($t_{\\rm age}\\lap 10^4$ yr) the orientation angle $\\alpha$ may grow, whereas in older pulsars it becomes smaller. It is possible that changes in $\\alpha$ are responsible for $n_{\\rm obs}<3$ for the Crab, PSR B1509-58, PSR B0540-69. Since the data intervals used to determine these braking indices contained no glitches, to explain these values of $n_{\\rm obs}$ the alignment angle $\\alpha$ would have to grow {\\em between} glitches. The changes in $\\alpha$ inferred from the Crab's persistent offsets might produce measurable changes in the pulse profile. For example, a change of $\\alpha$ at a glitch could change the duration of the line-of-sight's traverse through the pulse emission cone. The magnitude of the associated change in the total pulsed flux is $\\sim|\\Delta\\alpha|/w_{1/2}$, where $w_{1/2}$ is the half-width of the emission cone. Taking $\\alpha=86^\\circ$, determined from the radio data (\\cite{crab_alpha}), $\\Delta\\alpha$ for the 1989 glitch (which exhibited the largest observed offset) is $\\simeq 3\\times 10^{-3}$ rad (see eq. \\ref{alphashift}). From these numbers we estimate the relative change in pulsed flux to be $\\sim 1$\\%. We estimate a somewhat smaller change, $\\sim 0.4$\\%, using the value $\\alpha=80^\\circ$ obtained by Yadigaroglu and Romani (1996) in their analysis of gamma ray and radio data. Such flux changes may be measurable in either x-ray or optical bands. Sekimoto \\etal\\ (1995), using GINGA, found an upper limit on the change in the pulsed x-ray emission (1-6 keV) across the 1989 glitch of $<1.6$\\%. Moreover, Jones, Smith and Nelson (1980), found that the pulsed optical emission varied by $\\sim 1$\\% over 7 yr." + }, + "9701/hep-ph9701276_arXiv.txt": { + "abstract": "\\noindent Modern cosmology has created a tight link between particle physics / field theory and a wealth of new observational data on the structure of the Universe. These lecture notes focus on some of the most important aspects concerning the connection between theory and observations. The lectures begin with an overview of some recent progress and problems in inflationary cosmology. In particular, a pedagogical discussion of the theory of reheating is presented. The second topic is a survey of the theory of cosmological perturbations, the cornerstone of modern cosmology. The focus is on the gauge-invariant classical and quantum theory of fluctuations. The third topic concerns the role of topological defects in cosmology. Reviews of the cosmic string theory of galaxy formation and of defect-mediated GUT and electroweak baryogenesis are given. ", + "introduction": "Most aspects of high energy physics beyond the standard model can only be tested by going to energies far greater than those which present accelerators can provide. Fortunately, the marriage between particle physics and cosmology has provided a way to ``experimentally\" test the new theories of fundamental forces. The key realization, discovered both in the context of the inflationary Universe scenario$^{\\cite{Guth}}$ and of topological defects models$^{\\cite{ZelVil}}$ is that physics of the very early Universe may explain the origin of structure in the Universe. It now appears that a rich set of data concerning the nonrandom distribution of matter on a wide range of cosmological scales, and on the anisotropies in the cosmic microwave background (CMB), may potentially be explained by high energy physics. In addition, studying the consequences of particle physics models in the context of cosmology may lead to severe constraints on new microscopic theories. Finally, particle physics and field theory may provide explanations of some deep cosmological puzzles, e.g. why the Universe at the present time appears so homogeneous, so close to being spatially flat, and why it contains the observed small net baryon to entropy ratio. In these lectures, I focus on three important aspects of modern cosmology. The first concerns some fundamental problems of inflationary cosmology. In particular, some recent progress in the understanding of ``reheating\" in inflation will be reviewed. The second topic is the classical and quantum theory of cosmological perturbations, the main tool of modern cosmology. A general relativistic and quantum mechanical analysis of the generation and evolution of linearized fluctuations is essential in order to be able to accurately calculate the amplitude of density perturbations and CMB anisotropies. As a third topic, I discuss the role of topological defects in baryogenesis and as possible sees for cosmological structure formation. The specific outline is as follows: \\begin{enumerate} \\item{} {\\bf Introduction and Outline} \\item{} {\\bf Lecture 1: Inflationary Universe: Progress and Problems} \\\\{2.A} Problems of Standard Cosmology \\\\{2.B} Inflationary Universe Scenario \\\\{2.C} Problems of Inflation \\\\{2.D} Inflation and Nonsingular Cosmology \\\\{2.E} Reheating in Inflationary Cosmology \\\\{2.F} Summary \\item{} {\\bf Lecture 2: Classical and Quantum Theory of Cosmological Perturbations} \\\\{3.A} Basic Issues \\\\{3.B} Newtonian Theory \\\\{3.C} Relativistic Theory: Classical Analysis \\\\{3.D} Relativistic Theory: Quantum Analysis \\\\{3.E} Summary \\item{} {\\bf Lecture 3: Topological Defects, Structure Formation and Baryogenesis} \\\\{4.A} Quantifying Data on Large-Scale Structure \\\\{4.B} Topological Defects \\\\{4.C} Formation of Defects in Cosmological Phase Transitions \\\\{4.D} Evolution of Strings and Scaling \\\\{4.E} Cosmic Strings and Structure Formation \\\\{4.F} Specific Predictions \\\\{4.G} Principles of Baryogenesis \\\\{4.H} GUT Baryogenesis and Topological Defects \\\\{4.I} Electroweak Baryogenesis and Topological Defects \\\\{4.J} Summary \\end{enumerate} Unless otherwise specified, units in which $\\hbar = c = k_B = 1$ will be used. Distances are expressed in Mpc (1pc $\\simeq$ 3.06 light years). Following the usual convention, $h$ indicates the expansion rate of the Universe in units of $100$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega = \\rho / \\rho_c$ is the ratio of the energy density $\\rho$ to the critical density $\\rho_c$ (the density which yields a spatially flat Universe), $G$ is Newton's constant and $m_{pl}$ is the Planck mass. ", + "conclusions": "" + }, + "9701/astro-ph9701025_arXiv.txt": { + "abstract": "We derive new effective radii and total magnitudes for 5 E and S0 galaxies in the \\leo\\ group from wide-field CCD images. These are used in conjunction with recent literature velocity data to construct the fundamental plane (FP) of the \\leo\\ group. The rms scatter that we find is only 6 \\% in distance. The zero point of this relation provides a calibration of the FP as a distance indicator and directly determines the angular diameter distance ratio between the \\leo\\ group and more distant clusters. In the language of Jerjen and Tammann (1993) we determine a cosmic velocity of the \\leo\\ group of $757\\pm 68 \\kms$ relative to the Coma cluster, or $796\\pm 57 \\kms$ relative to a frame of 9 clusters. Combining this velocity with the Cepheid distance to M96, a member of \\leo, we find the Hubble constant to be $H_0=67\\pm 8\\ho$ or $H_0=70\\pm7\\ho$ for each case. The distance we obtain for the Coma cluster itself ($108\\pm12$ Mpc) is in good agreement with a number of other recent estimates. ", + "introduction": "Despite the successful determination of Cepheid distances to several spiral galaxies in the Virgo cluster and elsewhere with the Hubble Space Telescope (HST), the dispute over the value of the Hubble constant is still not settled (Tanvir \\etal\\ 1995; Mould \\etal\\ 1995; Sandage \\etal\\ 1996). Part of the problem undoubtedly lies with the uncertain peculiar velocity and significant depth, particularly for the late-type galaxies, of the Virgo cluster. These uncertainties are removed if secondary indicators can be used to extend distance measurements directly to more remote and, ideally, more compact clusters. Moreover, due to the morphological segregation of early type galaxies towards the cores of clusters, secondary distance indicators based on E and S0 galaxies are preferred. In this \\paper\\ we calibrate the so-called fundamental plane (FP) distance indicator (Djorgovski \\& Davis 1987; Dressler \\etal\\ 1987) in the \\leo\\ group. The \\leo\\ group (Ferguson \\& Sandage 1990) is the nearest group of galaxies containing both bright spirals as well as early type galaxies. The center of the group is defined by the E1 galaxy NGC 3379 \\mbox{(= M105)} and the S0 galaxy NGC 3384 which form a close pair. Planetary nebula luminosity function (PNLF) and surface brightness fluctuations (SBF) observations (Ciardullo, Jacoby, \\& Tonry 1993) indicate that the E5 galaxy NGC 3377 is a member of the group, whilst the radial velocity and projected distance of the S0 galaxy NGC 3412 from the core of the group indicate that it too is a likely member (Garcia 1993; Tanvir 1996). Finally, the S0/a galaxy NGC 3489 is also a possible early-type member of \\leo, however, given its distance from the group center ($\\sim3^{\\circ}$) and its comparatively ``late-type'' properties, we decided {\\it a priori} that it should not be included in the construction of the fundamental plane (see \\S 3). The fundamental plane is an excellent distance indicator relating distances between E and S0 galaxies particularly in clusters and groups. The most recent paper on the subject studies the FP for 230 E and S0 galaxies in 11 clusters and groups (J\\o rgensen, Franx \\& Kj\\ae rgaard 1996; hereafter JFK). For the present work, the important result of that study is the demonstration that a single global FP yields unbiased distance estimates to early-type galaxies in very different environments with a formal intrinsic scatter of 14 \\% in distance per galaxy. Thus with four early-type galaxies in the \\leo\\ group we may obtain the zero-point of the relationship to 7\\%, with the additional uncertainty of the absolute distance to the group. The distance to \\leo\\ has recently been measured by Cepheids in the spiral galaxies M96 (Tanvir \\etal\\ 1995) and M95 (Graham \\etal\\ 1997), by SBF (Tonry \\etal\\ 1997; Sodemann and Thomsen 1995; 1996), by the tip of the red-giant branch (TRGB) in NGC 3379 (Sakai \\etal\\ 1997) and PNLF in several group members (Ciardullo, Jacoby and Ford 1989; Feldmeir, Ciardullo, \\& Jacoby 1997). Thus the \\leo\\ group is an important stepping stone in the determination of the Hubble constant. ", + "conclusions": "The main result of this \\paper\\ is the construction of the fundamental plane (FP) of the \\leo\\ group. The small observed scatter (6 \\%) about this relation means that the result is insensitive to the particular choice of and weighting of the galaxy sample. We have adopted the form of FP determined by JFK and hence find the relative distance between Coma and \\leo\\ to be $9.5\\pm0.7$. In addition, using the distance to M96 in \\leo\\ ($11.3\\pm0.9$ Mpc) yields a direct and accurate measurement of the Coma distance of $108\\pm12$ Mpc. Finally, using the recession velocity of Coma ($7200 \\pm400\\kms$), we estimate the Hubble constant to be $H_0=67\\pm8\\ho$, or alternatively, if we consider the full cluster sample of JFK, with a median redshift $\\sim5000\\kms$, we find $H_0=70\\pm7\\ho$." + }, + "9701/astro-ph9701163_arXiv.txt": { + "abstract": "One way of recovering information about the initial conditions of the Universe is by measuring features of the cosmological density field which are preserved during gravitational evolution and galaxy formation. In this paper we study the total number density of peaks in a (galaxy) point distribution smoothed with a filter, evaluating its usefulness as a means of inferring the shape of the initial (matter) power spectrum. We find that in numerical simulations which start from Gaussian initial conditions, the peak density follows well that predicted by the theory of Gaussian density fields, even on scales where the clustering is mildly non-linear. For smaller filter scales, $r \\simlt 4-6 \\hmpc$, we see evidence of merging as the peak density decreases with time. On larger scales, the peak density is independent of time. One might also expect it to be fairly robust with respect to variations in biasing, i.e. the way galaxies trace mass fluctuations. We find that this is the case when we apply various biasing prescriptions to the matter distribution in simulations. If the initial conditions are Gaussian, it is possible to use the peak density measured from the evolved field to reconstruct the shape of the {\\it initial} power spectrum. We describe a stable method for doing this and apply it to several biased and unbiased non-linear simulations. We are able to recover the slope of the linear matter power spectrum on scales $k \\simlt 0.4 \\hmpc^{-1}$. The reconstruction has the advantage of being independent of the cosmological parameters ($\\Omega$, $\\Lambda$, $H_0$) and of the clustering normalisation ($\\sigma_8$). The peak density and reconstructed power spectrum slope therefore promise to be powerful discriminators between popular cosmological scenarios. ", + "introduction": "Attempts to reconstruct the initial conditions from which large-scale structure in the Universe grew have been carried out using several different methods. One avenue of research involves attempting to ``run gravity backwards'' from the present day galaxy distribution. This has been done with dynamical schemes (eg. Peebles 1989, Nusser \\& Dekel 1992, Croft \\& Gazta\\~{n}aga 1997) and also by applying a one to one mapping between the final smoothed density and the initial density (assumed to have a Gaussian p.d.f.) (Weinberg 1992). A different approach attempts to recover statistical measures of the initial conditions (eg. the two-point correlation function or the power spectrum) from knowledge of how gravitational evolution has affected these statistics. For example, formulae have been proposed which map the correlation function measured at the present day onto the initial correlation function (Hamilton $\\etal$ 1991). The same is true for the power spectrum (Peacock and Dodds 1994, Jain, Mo \\& White 1995 , Baugh \\& Gazta\\~{n}aga 1996). In this paper we will show how the initial power spectrum shape can be inferred from measurements of the space density of peaks in the galaxy distribution, even on scales where significant non-linear evolution has taken place in the underlying mass. Peaks in the initial density distribution have been thought to be potential sites for the formation of galaxies and clusters of galaxies. Two seminal papers by Kaiser (1984) and Bardeen $\\etal$ (1986) in which the details of the theory of Gaussian random fields relevant to these problems were studied have been very influential in the study of galaxy and large-scale structure formation. In the present paper we apply these results to the density field smoothed on larger scales, in order to see how peak theory describes the properties of an evolved density field. We are particularly interested in seeing whether the distribution of peaks enables us to recover some information about the initial conditions, and also in its potential use as a discriminator between cosmological models. When applying peak theory to the problem of galaxy or cluster formation, the assumption is usually made that these objects have formed from the gravitational collapse of matter around a high peak. Some smoothing scale, somewhat arbitrarily chosen, is associated with the object to be formed. The validity of these assumptions has been tested on the small scales relevant to the formation of individual objects by Park (1991), and Katz $\\etal$ (1993) amongst others. The conclusions seem to be that at least for the formation of matter haloes, there is only qualitative agreement between the predictions and the results of numerical simulations. Indeed in view of the extensive merging of matter clumps predicted by presently favoured hierarchical models of structure formation (see e.g. Press $\\&$ Schecter 1974, Lacey $\\&$ Cole 1993) it is reasonable to expect that evolution of the density field on small scales will tend to disrupt the predictions of peak theory. When the matter distribution is smoothed on somewhat larger scales however, for example in the estimation of the genus (e.g. Melott, Weinberg $\\&$ Gott 1988), the topology and some aspects of large scale features of the density field seem to be little affected by gravitational evolution. In particular, the rank order of density contrasts measured at points in the field is approximately conserved. This fact has been used by Weinberg (1992) in the reconstruction method mentioned above. This reconstruction method also appears to work well when tested on simulations where the matter distribution is biased using some ad hoc galaxy formation model. In this paper we will use numerical simulations to test whether this insensitivity to gravitational evolution affects the mean number of peaks and if so on what scales it begins to break down. We are also interested in reconstructing the power spectrum shape from the peak density as one specifies the other entirely in Gaussian models. The outline of the paper is as follows: In Section 2 we will briefly introduce the concept of the number density of peaks in Gaussian random fields and in $N$-body simulations. We will describe the simulations and our method of peak selection before testing peak theory against the evolved density both of mass and of `galaxies' identified in the models. In Section 3 (and Appendix A) we will describe a stable method of recovering the power spectrum slope from the peak density. We will test it first on analytic power spectra and then on the simulation results. In Section 4 we discuss our results and present some conclusions. ", + "conclusions": "Although a local transformation of a density field can change the heights of peaks, it should not affect the peak number density. Thus, as far as the effects of both gravity and biasing are approximately local one should expect the galaxy peak density to be a preserved quantity. Gravity is local on linear scales and there has not been enough cosmic time for gravity to affect (or become non-linear on) scales much larger than $\\sim 8 \\mpc$. For similar reasons it also seems difficult for biasing to be significantly non-local (see also Gazta\\~{n}aga \\& Frieman 1994 and references therein). Of course, it may also be possible to have non-local transformations which preserve the peak number density. Even without attempting to reconstruct the power spectrum shape from it, the space density of peaks appears to be an interesting statistic. At the present time, the two cosmological models which are the most favoured alternatives to $\\Gamma=0.5$ CDM are Low density CDM and MDM, which we have seen have very differently shaped linear power spectra and different peak space densities. The similarity of their shapes on quasilinear scales, and the uncertainties from biasing on smaller scales mean that galaxy clustering observations generally favour neither one or the other. We propose that observational measurements of the peak space density would constitute a simple means of telling them apart. Attempts have already been made to constrain models based on linear power spectra reconstructed from the non-linear observations using mapping formulae calibrated against $N$-body simulations (as mentioned in Section 1). Unfortunately, two different analyses have come up with different conclusions, each favouring a different model, Low density CDM for Peacock \\& Dodds (1994) and MDM for Baugh \\& Gaztanaga (1996). This may be because of the sensitivity of the non-linear mapping formulae to cosmological parameters such as $\\Omega$, or the fact that they work less well with steep power spectra. The reconstruction from the peak space density should not be sensitive to the values of $\\Omega$ and $\\Lambda$, for example, and should be more insensitive to variations in galaxy biasing. It should be borne in mind, though, that the work in this paper is aimed at the quasilinear regime and that the non-linear mapping formulae can, in principle, be used as a tool to study much smaller scales (although the uncertainties due to biasing could be large). If one assumes that the initial conditions were Gaussian, we have shown that it is possible to reconstruct the shape of P(k) on interesting scales from the peak density. The method appears to work best for the models with the steeper slopes on small scales, MDM and $\\Gamma=0.2$ CDM. This is just as well, as these models are more favoured by galaxy clustering data. The effects of sparse sampling, boundary conditions and redshift distortions must be studied in conjunction with an application to real data. As with the genus statistic, a large contiguous volume is necessary in order to estimate $n_{pk}$ reliably. With the next generation of large redshift surveys (Colless 1997, Gunn \\& Weinberg 1995) it should be possible to measure $n_{pk}$ with comparable accuracy to the simulations we have presented here, and probably with even better accuracy on large scales. In conclusion, we have shown that the space density of peaks in a density field smoothed with a Gaussian filter is independent of time, following the linear theory prediction, even on scales where the clustering is mildly non-linear. We have also shown that the peak space density in itself can be used to distinguish between cosmological models, including the currently popular MDM and Low density CDM scenarios. The peak density could also be used to compare with non-Gaussian models, if predictions become available. We have developed a simple, stable method of recovering the power spectrum from the space density of peaks and demonstrated that the method works using analytic power spectra and the simulation results. An application of the method to a contiguous, densely sampled galaxy survey should yield a useful quantity, the shape of the linear power spectrum on scales $k \\simlt 0.4 \\hmpc^{-1}$." + }, + "9701/gr-qc9701039_arXiv.txt": { + "abstract": "We estimate the signal-to-noise ratios (SNRs) that one would expect to measure from coalescing binary black hole (BBH) systems for the following three broadband gravitational-wave observatories: initial and advanced ground-based interferometers (LIGO/VIRGO) and space-based interferometers (LISA). We focus particularly on the highly relativistic and nonlinear {\\it merger} portion of the gravitational-wave signal, which comes after the adiabatic {\\it inspiral} portion and before the {\\it ringdown} portion due to the quasinormal ringing of the final Kerr black hole. Ground-based interferometers can do moderate SNR (a few tens), moderate accuracy studies of the dynamics of merging black holes in the mass range $({\\rm a~few}) M_\\odot$ to $\\sim 2000 M_\\odot$. LISA, by contrast, can do high SNR (a few $\\times\\,10^4$), high-accuracy studies of BBH systems in the mass range $10^5 M_\\odot\\alt (1+z) M \\alt 10^8 M_\\odot$, where $z$ is the binaries' cosmological redshift. Our estimated SNRs suggest that coalescing black holes might well be the first sources detected by the LIGO/VIRGO network of ground-based interferometers. Because of their larger masses, they can be seen out to much greater distances (up to $\\sim 250 \\, {\\rm Mpc}$ for $M \\alt 50 M_\\odot$ for the initial LIGO interferometers) than coalescing neutron star binaries (heretofore regarded as the ``bread and butter'' workhorse source for LIGO/VIRGO, visible to $\\sim 30 \\, {\\rm Mpc}$ by the initial LIGO interferometers). Low-mass BBHs ($M \\alt 30 M_\\odot$ for the first LIGO interferometers; $M \\alt 80 M_\\odot$ for the advanced; $(1+z) M \\alt 3 \\times 10^6 M_\\odot$ for LISA) are best searched for via their well-understood inspiral waves; more massive BBHs must be searched for via their more poorly understood merger waves and/or their well-understood ringdown waves. A search for low-mass BBHs based on the inspiral waves, at a sensitivity level roughly half way between the first LIGO interferometers and the advanced LIGO interferometers, should be capable of finding BBHs out to $\\sim 1 \\, {\\rm Gpc}$. A search for massive BBHs based on the ringdown waves can be performed using the method of matched filters. If one wants the reduction in the event rate due to the discreteness of the template family to be no more than $10\\%$, then the number of independent templates needed in the search will only be about $6000$ or less. Such a search with the first LIGO interferometers should be capable of finding BBHs in the mass range from about $100M_\\odot$ to about $700 M_\\odot$ out to $\\sim 200 \\, {\\rm Mpc}$; with advanced LIGO interferometers, from about $200 M_\\odot$ to about $3000 M_\\odot$ out to $z \\sim 1$; and with LISA, BBHs with $10^6 M_\\odot \\alt (1+z) M \\alt 3 \\times 10^8 M_\\odot$ should be visible out to $z \\agt 100$. The effectiveness of a search based on the merger waves will depend on how much one has learned about the merger waveforms from numerical relativity simulations. With only a knowledge of the merger waves' range of frequency bands and range of temporal durations, a search based on merger waves can be performed using a nonlinear filtering search algorithm. Such a search should increase the number of discovered BBHs by a factor of roughly $10$ over those found from the inspiral and ringdown waves. On the other hand, a full set of merger templates based on numerical relativity simulations could further increase the number of discovered BBHs by a additional factor of up to about $4$. ", + "introduction": "\\label{intro} \\subsection{Coalescences of black hole binaries} It has long been recognized that coalescences of binary systems of two black holes could be an important source of gravitational waves \\cite{300yrs,schutzreview}, both for the ground based interferometric detectors LIGO \\cite{ligoscience} and VIRGO \\cite{virgo} currently under construction, and also for the possible future space-based interferometer LISA \\cite{cornerstone,lisa,LISAreport}. The orbits of binary black holes (BBHs) gradually decay from energy and angular momentum loss to gravitational radiation. Eventually, the holes coalesce to form a final black hole. For gravitational radiation reaction to successfully drive the binary to merge in less than a Hubble time, the initial orbital period must be $\\alt 0.3 \\, {\\rm days} \\, (M/M_\\odot)^{5/8}$, where $M$ is the total mass of the binary; thus the critical orbital period is of order days for solar mass black holes, and of order years to hundreds of years for supermassive black holes ($10^6 M_\\odot \\alt M \\alt 10^9 M_\\odot$). The process of coalescence can be divided up into three more or less distinct phases: \\begin{itemize} \\item An adiabatic {\\it inspiral}, during which the gravitational radiation reaction timescale is much longer than the orbital period. The inspiral ends when the binary orbit becomes relativistically dynamically unstable at an orbital separation of $r \\sim 6 M$ (in units where $G=c=1$) \\cite{kww,cook}. The gravitational waves from the inspiral carry encoded within them the masses and spins of the two black holes, some of the orbital elements of the binary and the distance to the binary \\cite{300yrs,Kipreview}. \\item Towards the end of the inspiral, the black holes encounter the dynamical instability and make a gradual transition from a radiation-reaction driven inspiral to a freely-falling plunge \\cite{kww,KipAmos,laiwiseman}, after which, even if the radiation reaction could be turned off, the black holes would still merge. We will call the subsequent plunge and violent collision the {\\it merger} phase. Gravitational waves from the merger could be rich with information about the dynamics of relativistic gravity in a highly nonlinear, highly dynamic regime about which we have a poor theoretical understanding today. \\item As the final black hole settles down to a stationary Kerr state, the nonlinear dynamics of the merger gradually become more and more describable as oscillations of this final black hole's quasinormal modes \\cite{pressteukolsky,chandradetweiler}. The corresponding emitted gravitational waves consist of a superposition of exponentially damped sinusoids. We will call the phase of the coalescence for which the emitted gravitational-wave signal is dominated by the strongest $l=m=2$ quasinormal mode signal the {\\it ringdown} phase. The waves from the ringdown carry information about the mass and spin of the final black hole \\cite{echeverria,finnmeasure}. (Note that, for want of a better terminology, throughout this paper we consistently use {\\it coalescence} to refer to the entire process of inspiral, merger and ringdown, and reserve the word merger for the phase intermediate between inspiral and ringdown.) \\end{itemize} In this paper we focus primarily on BBHs in which the masses of the two black holes are approximately the same, although we do also consider sources where one BH is much smaller than the other. We consider three different classes of BBHs: \\medskip (i) {\\it Solar mass} black hole binaries, {\\it i.e.}, binaries that are formed from massive main-sequence progenitor binary stellar systems. Such BBHs are expected to have total masses in the range $10 M_\\odot \\alt M \\alt 50 M_\\odot$, but not much larger than this. The rate of coalescence of solar-mass BBHs in the Universe is not very well known. However, theory suggests that most of the BBH progenitor systems may not disrupt during the stellar collapses that produce the black holes, so that their coalescence rate could be about the same as the birth rate for their progenitors, about $1/100,000$ years in our Galaxy, or several per year within a distance of $200 \\, {\\rm Mpc}$ \\cite{narayan,phinney,heuvel,tutukov,yamaoka}. Note that this coalescence rate is roughly the same as the expected event rate for what has traditionally been regarded as the most promising source for ground based interferometers, coalescences of neutron star---neutron star (NS-NS) binaries \\cite{ligoscience,Kipreview}. The expected rate of NS-NS coalescences is more firmly known, however, since it is based on extrapolations from detected progenitor NS-NS systems \\cite{narayan,phinney,heuvel}. \\medskip (ii) {\\it Intermediate mass} black hole binaries, with total masses in the range $50 M_\\odot \\alt M \\alt ({\\rm a~few}) \\times 10^3 M_\\odot$. In contrast to the cases of solar mass black holes and supermassive black holes (discussed below), there is little direct observational evidence for the existence of black holes in this mass range. Although there have been suggestions that the globular cluster M15 harbors a black hole of mass $\\sim 10^3 M_\\odot$ \\cite{postnov1}, theoretical modeling combined with the most recent HST observations neither confirm nor rule out this possibility \\cite{M15}. Despite the lack of observational evidence, it is plausible that black holes in this mass range are formed in the cores of globular clusters, or in galactic nuclei in the process of formation of a supermassive black hole \\cite{Rees}. Simulations by Quinlan and Shapiro suggest that black holes with $M \\sim 100 M_\\odot - 1000 M_\\odot$ could be formed in the evolution of dense stellar clusters of main sequence stars in galactic nuclei \\cite{QuinlanShapiro}, and that coalescences of binaries of such black holes could be possible en route to the formation of a supermassive black hole. We are prompted to consider these intermediate mass BBHs by the following consideration: even if the coalescence rate of intermediate mass BBHs is $\\sim 10^{-4}$ that of NS-NS binaries (which is thought to be $\\sim 10^{-5} \\, {\\rm yr}^{-1}$ in our Galaxy as discussed above), these sources would still be seen more often than NS-NS binaries by LIGO's initial and advanced interferometers, and thus could be the first detected type of source. (See Sec.~\\ref{results} for further details.) \\medskip (iii) {\\it Supermassive} black holes binaries: There is a variety of strong circumstantial evidence that supermassive black holes (SMBHs) in the mass range $10^6 - 10^9 M_\\odot$ are present in quasars and active galactic nuclei, and also that $\\sim 25\\%$ -- $ 50\\%$ of nearby massive spiral and elliptical galaxies harbor quiescent SMBHs \\cite{LISAreport,BlandfordRees}. See Ref.~\\cite{Rees} for a review of this evidence. One of the main scientific goals of the LISA project is to detect and monitor various processes involving SMBHs, such as the capture of compact stars \\cite{schutzreview,LISAreport,Kipreview,HilsBender,Ericnew}, and their formation \\cite{schutzreview,LISAreport}. In particular, the coalescences of SMBH binaries that are formed in Galaxy mergers, in which the individual SMBHs are driven together by dynamical friction and gas accretion until gravitational radiation reaction takes over \\cite{Blandford}, have often been suggested as a promising source for space-based interferometers \\cite{300yrs,schutzreview,LISAreport,Kipreview,Wahlquist,Haehnelt}. Such coalescences would be detectable throughout the observable universe with large signal to noise ratios \\cite{LISAreport,Kipreview}. There is also observational evidence for SMBH binaries: wiggles in the radio jet of QSO 1928+738 have been attributed to the orbital motion of a SMBH binary \\cite{Roos}, as have time variations in quasar luminosities \\cite{sillanpana} and in emission line redshifts \\cite{gaskell}. The overall event rate is uncertain, but could be large ($\\agt 1/{\\rm yr}$), especially if the hierarchical scenario for structure formation is correct \\cite{Haehnelt}. \\subsection{Status of theoretical calculations of the gravitational-wave signal} \\label{calcstatus} Detailed theoretical understanding and predictions of the gravitational waveforms $h_+(t)$ and $h_\\times(t)$ produced in BBH coalescences will facilitate both the detection of the gravitational-wave signal, and the extraction of its useful information. In situations where a complete family of theoretical template waveforms is available, it will be possible to use the procedure of Wiener optimal filtering to search the interferometer data streams and to detect the gravitational-wave signal \\cite{300yrs,schutz1}. The resulting signal-to-noise ratios (SNRs) can be larger than those obtainable without theoretical templates by a substantial factor; see Sec.~\\ref{derivesnrformulagen}. Thus, while it is possible to detect the various phases of BBH coalescences without theoretical templates, such templates can greatly increase the effective range of the interferometers and the event detection rate. (Accurate theoretical templates are also essential for extracting the maximum amount of information from detected signals \\cite{paperII}.) Such theoretical template waveforms are available for the inspiral and ringdown phases of the coalescence, but not yet for the merger phase, as we now discuss. For the inspiral phase of the coalescence, the gravitational waves and orbital evolution can be described reasonably well using the post-Newtonian approximation to general relativity. To date, inspiral waveforms have been calculated up to post-2.5-Newtonian order \\cite{Blanchet}, and the prospects look good for obtaining waveforms up to post-3.5-Newtonian order \\cite{Blanchet1,ordercount}. Post-Newtonian templates will be fairly accurate over most of the inspiral, the most important error being a cumulative phase lag \\cite{prl,ericaccuracy}. This cumulative phase lag will not be important for searches for inspiral waves; template phasing error will be largely compensated for by systematic errors in best-fit values of the binary's parameters, and the signals will still be found \\cite{prl,bala1,cutlerflan2,searchtemplates}. By contrast, template inaccuracies will be significant when one attempts to extract from the data the binary's parameters. In particular, post-Newtonian templates' errors start to become very significant around an orbital separation of $r \\sim 12 M$ \\cite{PNbreakdown}, well before the end of the inspiral at the dynamical orbital instability ($r \\sim 6 M$). Templates for the phase of the inspiral between roughly $12 M$ and $6 M$ will most likely have to be calculated using methods other than the post-Newtonian approximation. Moreover, the methods of full blown numerical relativity cannot be applied to this ``Intermediate Binary Black Hole'' (IBBH) phase, since the total time taken to evolve from $12 M$ to $6 M$ is about $1500 M$, too long for supercomputer simulations to even contemplate evolving. Alternative analytic and numerical methods for calculating gravitational waveforms from the IBBH portion of BBH inspirals, based on the adiabatic approximation, are under development \\cite{IBBHworkshop}; it is likely that such alternative methods will be successfully developed and implemented before the gravitational-wave detectors begin their measurements \\cite{IBBH1}. Waveforms from the dynamic, complicated merger phase of the coalescence can only be obtained from numerical relativity. Unlike mergers of neutron star binaries, BBH mergers are particularly clean in the sense that there is no microphysics or hydrodynamics to complicate simulations of the evolution, and external perturbations are negligible: the entire merger can be described as a solution to the vacuum Einstein equation \\cite{notsoclean}. Finding that solution is not a particularly easy task: a major computational effort to evolve the vacuum Einstein equation for BBH mergers using massive computational resources is currently underway, funded by the National Science Foundation's Grand Challenge program \\cite{grandchall,samreview}. The ringdown phase of the coalescence can be accurately described using perturbation theory on the Kerr spacetime background \\cite{chandra}. The gravitational waveforms from this phase are well understood, being just exponentially damped sinusoids. Thus, Wiener optimal filtering is feasible for searches for ringdown waves. \\subsection{Purpose of this paper} \\label{whyarewehere} The principal purpose of this paper is to estimate, in more detail than has been done previously, the prospects for measuring gravitational waves from the three different phases of coalescence events (inspiral, merger and ringdown), for various different detectors, and for a wide range of BBH masses. We estimate in each case the distances to which the different types of source can be seen by calculating expected SNRs. In particular, we determine for each BBH mass and each detector whether a coalescence event is most effectively detected by searching for the inspiral portion of the signal, or the merger portion, or the ringdown portion. We also determine how much the availability of theoretical template waveforms for the merger phase could increase the event detection rate. Previous estimates of SNRs for ground-based interferometers have focused on the inspiral \\cite{300yrs,prl} and ringdown \\cite{echeverria,finnmeasure} phases, and also focused on solar-mass BBHs. For space-based interferometers, previous estimates of SNRs from the merger phase \\cite{LISAreport,Kipreview} were restricted to specific mass values and did not consider the ringdown portion of the signal. In a companion paper, we discuss in detail the useful information carried by the three phases of the gravitational-wave signal, and methods and prospects for extracting this information both with and without templates for the merger phase \\cite{paperII}. In the following subsection we describe our calculations and summarize the assumptions underlying our estimated signal-to-noise ratios. In Sec.~\\ref{results} we summarize our main results and conclusions, and an outline of the paper is given in Sec.~\\ref{map}. \\subsection{Estimating the signal-to-noise ratios: method and assumptions} \\label{assumptionsec} We calculate SNRs for three different types of interferometer: initial and advanced ground-based interferometers (LIGO/VIRGO), and the proposed space-based interferometer LISA. The noise spectra of the initial and advanced ground-based interferometers we took from Ref.~\\cite{ligoscience}, and that for LISA we obtained from Ref.~\\cite{LISAreport}. Our approximate versions of these noise spectra are given in Eqs.~(\\ref{noisespec}) -- (\\ref{noisespecLISA}), and are illustrated in Figs.~\\ref{hcharfig} -- \\ref{hcharfig2} in Sec.~\\ref{egs}. We consider the following three different signal-detection methods: \\medskip (i) {\\it Matched filtering searches:} For those phases of the coalescence for which a complete set of theoretical templates will be available (the inspiral, the ringdown, and possibly the merger), the method of matched filtering or optimal filtering can be used to search for the waves \\cite{300yrs,schutz1,helstrom,wainandzub,cutlerflan1}. The theory of matched filtering is briefly sketched in Sec.~\\ref{derivesnrformula1}. For any source of waves, the SNR, $\\rho$, obtained from matched filtering is related to the gravitational waveform $h(t)$ measured by the interferometer and to the spectral density $S_h(f)$ of the strain noise in the interferometer via \\cite{twosided} \\begin{equation} \\label{snr0} \\rho^2 = 4 \\int_0^\\infty {| {\\tilde h}(f)|^2 \\over S_h(f)} df, \\end{equation} where ${\\tilde h}(f)$ is the Fourier transform of $h(t)$ defined by Eq.~(\\ref{fourier}). The SNR (\\ref{snr0}) depends, through the waveform $h(t)$, on the orientation and position of the source relative to the interferometer. In Sec.~\\ref{derivesnrformula} we show that if we perform an rms average over source orientations and positions (at a fixed distance), the rms SNR thus obtained depends only on the energy spectrum $dE/df$ carried off from the source by the gravitational waves. The resulting relationship between the waves' energy spectrum and the rms angle-averaged SNR forms the basis for most of our calculations. It is given by [{\\it cf}.~Eq.~(\\ref{snraveraged})] \\begin{equation} \\langle \\rho^2 \\rangle = {2 (1 + z)^2 \\over 5 \\pi^2 D(z)^2} \\, \\int_0^\\infty df \\, {1 \\over f^2 S_h(f)} \\, {d E \\over d f}[(1 + z)f], \\label{snraveraged0} \\end{equation} where $z$ is the source's cosmological redshift and $D(z)$ is the luminosity distance from the source. In order for a signal to be detected, the waves' measured SNR must be larger than the detection threshold \\begin{equation} \\rho_{\\rm threshold} \\approx \\sqrt{ 2 \\ln( {\\cal N}_{\\rm start-times}) + 2 \\ln( {\\cal N}_{\\rm shapes})}; \\label{MFthreshold} \\end{equation} see, for example, Ref.~\\cite{prl} and also Sec.~\\ref{nonlinear}. Here ${\\cal N}_{\\rm start-times}$ is the number of independent starting times of the gravitational wave signal that are searched for in the data set, determined by the total duration of the data set (of order one year) and the sampling time. The quantity ${\\cal N}_{\\rm shapes}$ is the total number of statistically independent waveform shapes in the set of signals that one is searching for \\cite{importantnote}. \\medskip (ii) {\\it Band-pass filtering searches:} For the merger phase, a complete set of theoretical templates may not be available. It therefore may not be possible to use matched filtering, and other search methods will need to be employed. Band-pass filtering, followed by setting a detection threshold in the time domain, is a simple method of searching an interferometer data stream for bursts of unknown form \\cite{schutz1}. In Sec.~\\ref{derivesnrformula1} we derive an approximate relation between the SNR obtainable from band-pass filtering, and the SNR (\\ref{snr0}) obtainable from matched filtering, for any burst of waves, namely \\begin{equation} \\left({S \\over N} \\right)_{\\rm band-pass} \\approx {1 \\over \\sqrt{2 T \\Delta f}} \\, \\left({S \\over N} \\right)_{\\rm optimal}. \\label{bpa} \\end{equation} Here $T$ is the duration of the burst and $\\Delta f$ is the bandwidth of the band-pass filter [{\\it cf}. Eq.~(\\ref{simp1})]. The quantity $2 T \\Delta f$ is the dimension of the linear space of signals being searched for, which is roughly the same as the ``number of cycles'' of the gravitational waveform. We use this formula in Sec.~\\ref{detectmerger} to estimate band-pass-filter search SNRs for the merger waves from BBH coalescences, by inserting on the right hand side the rms angle-averaged matched-filter search SNR (\\ref{snraveraged0}), and by making estimates of $T$ and $\\Delta f$. \\medskip (iii) {\\it Noise-monitoring, nonlinear filtering searches:} The traditional view has been that the SNR (\\ref{bpa}) is about the best that can be achieved in the absence of theoretical templates; that is, that the gain in SNR obtainable from matched filtering is approximately the square root of the number of cycles in the gravitational wave signal. This view is based on the the assumption that the search method used in the absence of templates is band-pass filtering or something very similar. However, we suggest in Sec.~\\ref{nonlinear} an alternative search method, motivated by Bayesian analyses and incorporating nonlinear filtering, which performs much better than band-pass filtering and in some cases almost as well as matched filtering. In essence, one monitors the noise level in the interferometer in a certain frequency band, over short timescales, and looks for statistically significant changes. The noise level is estimated by calculating the quantity \\begin{equation} {1 \\over T} \\int_{-T/2}^{T/2} d \\tau \\, s(t + \\tau)^2, \\end{equation} where $T$ is the maximum expected duration of the signal, and $s(t)$ is a suitably pre-filtered version of the interferometer data stream. The efficiency of this noise-monitoring search method cannot usefully be described in terms of a signal-to-noise ratio, since the detection statistic is very non-Gaussian. Instead, its efficiency can be described in the following way. Let $\\rho$ denote the SNR that would be obtained if matched filtering were possible [Eq.~(\\ref{snr0})]. We use $\\rho$ as a convenient parameterization of the signal strength; as such, it is meaningful even in situations where matched filtering cannot be carried out. Then, a signal will be detected with high confidence using the noise-monitoring technique whenever $\\rho$ is larger than the threshold $\\rho_*$, where to a good approximation $\\rho_*$ satisfies \\FL \\begin{equation} \\rho_*^2 = 2 \\ln( {\\cal N}_{\\rm start-times}) + {\\cal N}_{\\rm bins} \\ln \\left( 1 + \\rho_*^2 / {\\cal N}_{\\rm bins} \\right). \\label{NLthreshold} \\end{equation} Here ${\\cal N}_{\\rm bins} = 2 T \\Delta f$ is the dimension of the signal space discussed above (or the number of independent frequency bins in the Fourier domain). The derivation of Eq. (\\ref{NLthreshold}) is given in Sec.~\\ref{nonlinear}. The relation (\\ref{snraveraged0}) forms the basis of our SNR calculations. We use the SNR thresholds (\\ref{MFthreshold}) and (\\ref{NLthreshold}) to deduce from the SNR values the detectability of the various parts of the gravitational-wave signal. To calculate the SNRs, we also need to specify the waves' energy spectra for the three different phases of the coalescence. As we now outline, for the inspiral and ringdown phases the waves' energy spectrum is essentially known, while for the merger phase we make an educated guess of $dE/df$. Sec.~\\ref{signalassumptions} gives more details. \\medskip {\\it Inspiral energy spectrum}: We use the leading order expression for $dE/df$ obtained using Newtonian gravity supplemented with the quadrupole formula \\cite{shapteuk} [Eq.~(\\ref{dEdfinspiral})]. Strictly speaking, this spectrum describes the SNR that would be achieved by searching for Newtonian, quadrupole waves using Newtonian, quadrupole templates. The actual SNR obtained when searching for a real, general-relativistic inspiral waveform using post-Newtonian templates should deviate from this by only a few tens of percent {\\cite{spectrumcomment}}. We terminate the spectrum at the frequency $f_{\\rm merge} = 0.02/M$ which is (roughly) the frequency of quadrupole waves emitted at the orbital dynamical instability at $r\\sim 6 M$ \\cite{kww}. For LISA, we assume that the measurement process lasts at most one year, and choose the frequency at which the inspiral spectrum starts accordingly. \\medskip {\\it Ringdown energy spectrum}: The spectrum that we use [Eq.~(\\ref{dEdfqnr})] is determined (up to its overall amplitude) by the characteristics of the $l=m=2$ quasi-normal ringing (QNR) mode of the final Kerr black hole. This mode is the most slowly damped of all QNR modes, so we expect it to dominate the last stages of the gravitational-wave emission. The QNR spectrum depends on three parameters: the quasi-normal modes' frequency $f_{\\rm qnr}$ and damping time $\\tau$, and the overall amplitude of the quasinormal mode signal. Equivalently, the three parameters can be taken to be the mass $M$ and dimensionless spin parameter $a$ of the final black hole (which determine $f_{\\rm qnr}$ and $\\tau$) and the total amount of energy radiated in the ringdown (which determines the overall amplitude). The spectrum is peaked at $f=f_{\\rm qnr}$ with width $\\Delta f \\sim 1/\\tau$. In our analyses, we (somewhat arbitrarily) assume that $a=0.98$. It seems likely that in many coalescences the spin of the final black hole will be close to maximal, since the total angular momentum of the binary at the end of the inspiral is $\\sim 0.9 M^2$ when the individual black holes are non-spinning \\cite{explainorb}, and the individual black hole spins can add to this. Exactly how close to maximally spinning the final black hole will be is a matter that probably will not be decided until supercomputer simulations---or observations---settle the issue. In any case, the ringdown SNR values that we obtain depend only weakly on our assumed value of $a$ [{\\it cf}.~Eq.~(\\ref{qnrsnrII})], for fixed total energy radiated in the ringdown. The overall amplitude of the ringdown signal depends upon one's delineation of where ``merger'' ends and ``ringdown'' begins, which is somewhat arbitrary. For equal-mass BBHs, we assume a value of the overall amplitude that corresponds to a total radiated energy in the ringdown of $0.03 M$---a radiation efficiency of $3\\%$. This number is based on a back-of-the-envelope, quadrupole-formula-based estimate of the QNR mode's amplitude when the distortion of the horizon of the black hole is of order unity ({\\it cf}.~Sec.~\\ref{ringdownphase}). Although this radiation efficiency seems rather high, there have been numerical simulations of the evolution of distorted, spinning black holes in which the ringdown waves carry away $\\agt 3\\%$ of the black hole's total mass energy \\cite{seidel}. For non equal-mass black holes, we assume that the total energy radiated is $F(\\mu/M) 0.03 M$, where $F(\\mu/M) = (4 \\mu/M)^2$ and $\\mu$ is the reduced mass of the binary. This function gives the correct results for the equal-mass case and also gives the correct scaling law in the regime $\\mu \\ll M$; for general mass ratios the scaling law is probably a good approximation. \\medskip {\\it Merger energy spectrum}: Realistic merger energy spectra will vary substantially from event to event (depending on the initial spins of the inspiraling black holes). Currently, we have very little concrete information about such spectra, pending supercomputer simulations of BBH mergers. In Sec.~\\ref{mergerphase} we describe various circumstantial pieces of evidence, culled from the literature, relevant to merger spectrum. Based on that evidence, we adopt the following crude model for equal-mass BBHs: we assume a flat spectrum $dE/df = {\\rm const}$ extending from the frequency $f_{\\rm merge}=0.02/M$ of quadrupole waves at the end of inspiral to the quasinormal ringing frequency $f_{\\rm qnr} =0.13/M$, with amplitude such that the total radiated energy in the merger is $10\\%$ of the total mass energy of the spacetime [Eq.~(\\ref{dEdfmerger})]. We outline in Sec.~\\ref{mergerphase} two different ``handwaving'' arguments which suggest that in favorable cases the merger radiation efficiency may be as high as our assumed value of $\\sim 10\\%$. One of these arguments, due originally to Smarr \\cite{smarrthesis} and explored by Detweiler \\cite{detweiler2}, is based on extrapolation of perturbation theory results; the other argument is based on angular momentum conservation. Our assumed radiation efficiencies of $3\\%$ and $10\\%$ for the ringdown and merger phases should perhaps be interpreted as reasonable upper bounds that could be achieved in favorable cases, rather than as best-guess estimates. We note that numerical simulations that have been performed to date (which are restricted to axisymmetric situations) generally yield lower radiation efficiencies than we have assumed \\cite{baker}; moreover, these axisymmetric simulations generally find that ringdown waves carry most of the radiated energy. In Sec.~\\ref{mergerphase} we argue that the radiated energy in the merger phase could be boosted by the lack of symmetry in generic black hole mergers, and especially by the individual black holes' spins (if these spins are large). For non-equal mass BBHs, we again simply reduce the energy spectrum by the factor $F(\\mu/M) = (4 \\mu/M)^2$, while the upper and lower frequencies $f_{\\rm merge}$ and $f_{\\rm qnr}$ are taken to be independent of $\\mu$. \\subsection{Signal-to-noise ratios: results and implications} \\label{results} By inserting our assumed energy spectra (\\ref{dEdfinspiral}), (\\ref{dEdfmerger}) and (\\ref{dEdfqnr}) into Eq.~(\\ref{snraveraged0}), we obtain optimal-filtering SNRs for the three different phases of BBH coalescences as a function of the redshifted total mass $(1+z) M$ of the binary. The results are shown in Figs.~\\ref{initialligosnr}, \\ref{advancedligosnr} and \\ref{lisasnr} and formulae summarizing the results are given in Appendix \\ref{appsnr}. Also in Sec.~\\ref{mergernofilters} we estimate that, for the merger waves, the number of independent frequency bins ${\\cal N}_{\\rm bins}$ which characterize the signal falls roughly in the range $10 \\alt {\\cal N}_{\\rm bins} \\alt 30$, and that a conservative upper-bound is $\\sim 60$. We use this upper bound in Sec.\\ \\ref{detectmerger} to estimate the SNR threshold (\\ref{NLthreshold}) for merger waves using noise-monitoring searches when templates are unavailable. We discuss the implications of these SNRs and SNR thresholds in", + "conclusions": "It seems quite likely that the gravitational waves from merging BBH systems will be detected by the ground-based interferometers that are now under construction. The initial LIGO interferometers will be able to detect low mass ($\\alt 30 M_\\odot$) coalescences of equal-mass BBHs out to about $200 \\, {\\rm Mpc}$ via their inspiral waves, and higher mass ($100 M_\\odot \\alt M \\alt 700 M_\\odot$) systems out to about $200 \\, {\\rm Mpc}$ via their ringdown waves. The advanced LIGO interferometers will be able to detect equal-mass BBH coalescences in the mass range $10 M_\\odot \\alt M \\alt 300 M_\\odot$ to $z \\sim 1/2$ via their inspiral waves, and higher mass ($200 M_\\odot \\alt M \\alt 3000 M_\\odot$ ) systems to $z \\sim 1$ via their ringdown waves. For non-equal mass BBHs, these distances will be reduced by a factor of $\\sim \\sqrt{4 \\mu/M}$ for inspiral signals and $\\sim 4 \\mu/M$ for ringdown signals. Searches for massive BBHs ($M \\agt 50 M_\\odot$ for LIGO/VIRGO) based on merger waves could increase the range of the interferometers by a additional factor of $\\sim 2$, without requiring detailed knowledge of the waveform shapes. It seems likely that BBH coalescences will be detected early in the gradual process of improvement from the first interferometers to the advanced interferometers, rather than later, and there is a strong possibility that they will be the first sources of gravitational radiation to be detected. Theoretical template waveforms obtained from numerical relativity supercomputer simulations will be crucial for analyzing the measured merger waves. A match of the detected waveform with a predicted waveform would be a triumph for the theory of general relativity and an absolutely unambiguous signature of the existence of black holes. A complete set of such theoretical templates would also aid the search for BBHs, but not by a large amount. The space-based interferometer LISA will be an extremely high precision instrument for studying the coalescences of supermassive BBHs. Coalescences with masses in the range $10^6 M_\\odot \\alt (1+z) M \\alt 10^9 M_\\odot$ should be detectable out to $z \\sim 10$ with very large SNRs ($\\agt 10^3$), via their merger and ringdown waves. Additionally, systems in the mass range $10^4 M_\\odot \\alt (1+z) M \\alt 3 \\times 10^7 M_\\odot$ should be detected to similar distances and with SNRs $\\agt 10^2$ via their inspiral waves." + }, + "9701/astro-ph9701232_arXiv.txt": { + "abstract": "We perform an extensive investigation of the sensitivity to non-vanishing $\\nut$ mass in a large water \\v{C}erenkov detector, developing an analysis method for neutrino events originated by a supernova explosion. This approach, based on directional considerations, provides informations almost undepending on the supernova model. We analyze several theoretical models from numerical simulations and phenomenological models based on $SN1987A$ data, and determine optimal values of the analysis parameters so as to reach the highest sensitivity to a non-vanishing $\\nut$ mass. The minimal detectable mass is generally just above the cosmologically interesting range, $m\\sim 100\\:\\ev$, in the case of a supernova explosion near the galactic center. For the case that no positive signal is obtained, observation of a neutrino burst with Super-Kamiokande will anyhow lower the present upper bound on $\\nut$ mass to few hundred $\\ev$. ", + "introduction": " ", + "conclusions": "" + }, + "9701/astro-ph9701142_arXiv.txt": { + "abstract": "We analyze several aspects of the recently noted neutron star collapse instability in close binary systems. We utilize (3+1) dimensional and spherical numerical general relativistic hydrodynamics to study the origin, evolution, and parametric sensitivity of this instability. We derive the modified conditions of hydrostatic equilibrium for the stars in the curved space of quasi-static orbits. We examine the sensitivity of the instability to the neutron star mass and equation of state. We also estimate limits to the possible interior heating and associated neutrino luminosity which could be generated as the stars gradually compress prior to collapse. We show that the radiative loss in neutrinos from this heating could exceed the power radiated in gravity waves for several hours prior to collapse. The possibility that the radiation neutrinos could produce gamma-ray (or other electromagnetic) burst phenomena is also discussed. ", + "introduction": "\\label{sec:level1} In recent numerical studies of the relativistic hydrodynamics of close neutron star binaries in three spatial dimensions (Wilson \\& Mathews 1995; Wilson, Mathews \\& Marronetti 1996, henceforth WMM96), it was noted that as the stars approach coalescence they appear to experience a collapse instability. For an appropriate equation of state, binary neutron stars might generally become black holes many seconds prior to merger. If correct, this effect could have a significant impact on the anticipated gravity wave signal from neutron star binaries near coalescence. Such premerger collapse might also be associated with heating, neutrino production, and electromagnetic bursts as the released gravitational energy from collapse is converted into thermal energy of the stars. Moreover, the numerical evidence that such an instability exists poses a number of new questions such as the sensitivity of the instability to the specific equation of state employed, or the intrinsic spin and masses of the stars. One would also like to understand the time history of the collapse and any associated electromagnetic or neutrino emission. In this paper we present some new three dimensional (3D) calculations which begin to examine these issues. Unfortunately, however, such relativistic hydrodynamic calculations in three spatial dimensions are computationally expensive. A complete systematic study of this instability in three spatial dimensions will be long in coming. In this paper, however, we show that in large part this effect can be replicated in terms of modified one-dimensional spherical relativistic hydrodynamics. We show that the relativistic effects of placing stars in a close binary can be approximated by adding a term involving an average Lorentz-like factor which increases the effective gravitational forces on the stars. The collapse observed in the three dimensional calculations can be understood in this one-dimensional framework and one can survey easily the sensitivity of this effect to parameters characterizing the binary and the neutron star equation of state. We can also follow the possible precollapse compression and heating of the neutron star material. This provides a framework in which to model the possible associated neutrino and electromagnetic signals such as gamma-ray bursts. We show that significant heating and neutrino emission is possible as the stars gradually compress before they reach the collapse instability. During the heating epoch the associated neutrino and electromagnetic radiative losses may dominate over the power loss from gravitational radiation. ", + "conclusions": "We have made a survey of the compression, heating, and collapse of neutron stars in close binaries. In particular, we have developed a schematic model to describe when the collapse instability may occur as a function of initial neutron star mass and the EOS. We have also analyzed the possible heating of the neutron star interiors as the stars approach the collapse instability. We find that the stars may obtain quite high thermal energy and neutrino luminosity in the final seconds before collapse. This could have significant implications both for gravity wave and neutrino astronomy as follows: \\subsection{Implication for Gravity Wave Detectors} The analysis here indicates that the radiative neutrino luminosity could exceed the gravity luminosity for hours prior to the collapse instability. If so, this could have a profound influence on the inferred gravity wave signal. The loss of orbital angular momentum due to neutrinos and electromagnetic radiation will be considerably greater than that of two cold stable neutron stars. The merger will occur on a shorter time scale and the gravity wave signal will be dominated by the dynamics of heating and thermal radiation and not the gravity wave amplitude up to the point of the instability. Once the collapse instability is reached, we estimate that the formation of one or two black holes will occur rather abruptly. After collapse, however, the system may not appear simply as two black holes in vacuum. As has been observed in supernova calculations for some time (cf. Mayle \\& Wilson 1988; Wilson \\& Mayle 1993) this much neutrino radiation is likely to ablate electron-positron pairs together with baryonic material from the surface of the stars. Baryons ejected in this wind are likely to be present after collapse and may interact with the orbiting objects. To some extent they will provide material to accrete onto the remaining members (neutron star or black holes) of the binary. They may also provide a damping medium which could accelerate the decay of the orbit. Thirdly, this hot wind material may provide a medium in which to anchor the magnetic field lines of the precollapse stars (cf. Wilson 1975; Ruffini \\& Wilson 1975; Damour et al. 1978). We speculate that these effects may serve to accelerate the merger of the two black holes. The interaction of the stars with this medium may affect the dynamics of the black hole inspiral unless the material is ejected with sufficiently high velocity. Clearly, this is an area which warrants further investigation. If our speculation is correct then the gravity wave signal becomes a probe of the EOS, hydrodynamics, and thermodynamics of the neutrons stars as they approach and pass through this collapse instability. \\subsection{Implications for Gamma-Ray Bursts} The possibility that gamma-ray bursts could be generated by neutrino emission from coalescing neutron stars has been speculated upon for some time. Recently, Janka \\& Ruffert (1996) have made post-Newtonian hydrodynamics calculations of neutron star mergers and included the neutrino emission therefrom. They find high luminosities, but the time scales are so short ($\\sim $ msec) that they conclude that it will be difficult to model gamma-ray bursts by neutron star mergers. This short time scale stems from the time scale for mergers. This difficulty is avoided, however, in our model in which the time scale is set by the gradual compression of the stars. We estimate similar luminosities, but in our model the neutrino luminosity endures for much longer times, thus rendering the possibility of a gamma-ray burst more plausible. We have shown that significant heating and associated neutrino luminosity is possible in the last seconds before the collapse instability. This poses some interesting possibilities for cosmological models of gamma ray bursts. The thermal emission of neutrinos provides an environment for the generation of an $e^+-e^-$ pair plasma by $\\nu \\bar \\nu$ annihilation around the stars. The neutrino emission is occurring in the deepening gravitational well of the two stars. Their interactions will be enhanced by the curved space around the neutron stars. Furthermore, the region between the stars may provide an environment for the build up of neutrino and matter flux and the production of a pair plasma as desired in some gamma-ray burst scenarios (e.g. Piran \\& Shemi 1993). In addition to the collapse-induced neutrino emission itself, the escaping neutrinos are likely to generate a neutrino heated baryon wind from the stars (Mayle \\& Wilson 1988; Wilson \\& Mayle 1993). Unlike in supernovae, the velocity of this wind can be be quite high, particularly later in the evolution as the neutrino luminosity grows. This later emission of the high velocity wind could interact with matter emitted previously producing shock heating in environments of relatively low optical depth far from the stars. The interactions themselves may contribute to the production of a pair plasma. As a preliminary test of this scenario we ran a calculation of neutron stars instantly heated such that the surface temperature was $\\sim 5 $ MeV. We then followed the neutrino and matter transport using the numerical supernova model of Mayle \\& Wilson (1993). We observed a blow off of the outer layer ($\\sim 10^{-5}$ M$_\\odot$) of the neutron star. This material was accelerated to a speed corresponding to a relativistic $\\gamma$-factor of $\\sim10$. One possibility is that this high speed matter interacting with magnetic fields and/or interstellar clouds might produce $\\gamma$-rays. Finally, we also note that after collapse, the previously ejected material will continue to experience heating either by accretion onto the black holes or by ram pressure from the orbiting stars. Once present, this plasma might become anchored to magnetic field lines around the precollapse stars (Ruffini \\& Wilson 1975; Damour et al. 1978). The interactions and magnetic recombination of these field lines could also contribute to heating and pair plasma production. All of these processes may be occurring in the background of the remaining orbiting binary system from times prior to collapse until the final merger to a single black hole. This orbit period may lead to an underlying millisecond substructure in associated burst signals possibly consistent with observations. Clearly, this is an area which also warrants more investigation. Work along this line is underway to explore such effects as a possible framework in which to model cosmological gamma-ray bursts." + }, + "9701/astro-ph9701156_arXiv.txt": { + "abstract": " ", + "introduction": "In this lecture we will try to address the \"frequently asked questions about fractals\" in the field of large scale galaxy distribution. This paper takes its origin from a very interesting discussion we had at this meeting. A lot of points were raised, and we try to make clear the fundamentals ones. For a more detailed discussion we refer the reader to \\cite{cp92} \\cite{bslmp94} for a basic introduction to this approach, and to \\cite{pmsl97} and \\cite{slmp97} for a review on the more recent results. In Sec.2, we briefly introduce the basics concepts of fractal geometry and the methods of correlation analysis, that are usually used in Statistical Mechanics. Moreover we present the results of our analysis in the case of real galaxy and cluster three dimensional samples. In Sec.3 we discuss the main points that have been raised during the \"Ringberg discussion\". Finally in Sec.4 we summarize our main conclusions. ", + "conclusions": "In summary our main points are: \\item The highly irregular galaxy distributions with large structures and voids strongly point to a new statistical approach in which the existence of a well defined average density is not assumed a priori and the possibility of non analytical properties should be addressed specifically. \\item The new approach for the study of galaxy correlations in all the available catalogues shows that their properties are actually compatible with each other and they are statistically valid samples. The severe discrepancies between different catalogues that have led various authors to consider these catalogues as {\\it not fair}, were due to the inappropriate methods of analysis. \\item The correct two point correlation analysis shows well defined fractal correlations up to the present observational limits, from 1 to $1000\\hmp$ with fractal dimension $D \\simeq 2$. Of course the statistical quality and solidity of the results is stronger up to $100 \\div 200 \\hmp$ and weaker for larger scales due to the limited data. It is remarkable, however, that at these larger scales one observes exactly the continuation of the correlation properties of the small and intermediate scales. \\item These new methods have been extended also to the analysis of the number counts and the angular catalogues which are shown to be fully compatible with the direct space correlation analysis. The new analysis of the number counts suggests that fractal correlations may extend also to scales larger that $1000\\hmp.$ \\item The inclusion of the galaxy luminosity (mass) leads to a distribution which is shown to have well defined multifractal properties. This leads to a new, important relation between the luminosity function and that galaxy correlations in space. \\item It is worth to notice Kerscher \\etal \\cite{ker97} presented at this meeting the morphological analysis of the IRAS 1.2 Jy by means of the Minkowski functional. Their conclusion that the scale of homogeneity is \"considerably larger than $200 \\hmp$\", is in complete agreement with ours. Moreover they have done a morphological characterization of structures that is complementary to the studies of the correlations properties presented in this lecture. Finally, we would like to stress also that these authors find again the \"apparent homogenization\" due to sparse sampling: the same kind of effect has been discussed here and in \\cite{slgmp96}. \\item Finally one should note that there are various {\\it indirect} arguments and always require an interpretation based on some assumptions. The most {\\it direct} evidence for the properties of galaxy distribution arises from the correct correlation analysis of the 3-d volume limited samples that has been the central point of our work. \\end{itemize}" + }, + "9701/astro-ph9701226_arXiv.txt": { + "abstract": "The damped \\lya absorbers (DLAs) in quasar spectra are believed to be the progenitors of present--day disk galaxies. We examine the probability for microlensing of background quasars by stars in their DLAs. Microlensing by an individual star should magnify the continuum but not the broad emission lines of the quasars. Consequently, the equivalent width distribution of microlensed quasars would be distorted. We model a representative spiral galaxy as a closed system composed of a bulge, a disk, and a halo, and evolve the mass fraction of stars in the disk based on the observed metallicity of DLAs at high redshifts. The microlensing signatures are stronger if the halo of the galaxy is made of Massive Compact Halo Objects (MACHOs). In this case, the distortion imprinted by microlensing on the equivalent width distribution of quasar emission lines can be detected with high significance in a sample of $\\sim 10$ DLAs with HI column densities $N\\ga 10^{21}~{\\rm cm^{-2}}$ and absorption redshifts $z_{\\rm abs}\\la 1$. About a tenth of all quasars with DLAs ($N\\ga 10^{20}~{\\rm cm^{-2}}$) might show excess variability on timescales shorter than five years. A search for these signals would complement microlensing searches in local galaxies and calibrate the MACHO mass fraction in galactic halos at high redshifts. ", + "introduction": "Damped Ly$\\alpha$ absorbers (DLAs) are thought to be the progenitors of present--day disk galaxies (Wolfe 1995). The abundance of heavy elements at low ionization stages in the absorbers and the velocity field traced by them is consistent with values expected for disk galaxies (Turnshek et al. 1989; Wolfe et al. 1993; Lu et al. 1993; Pettini et al. 1994; Lu \\& Wolfe 1994). Recent HIRES observations on the Keck telescope indicate that the weak, low--ionization, metal absorption lines in these systems often show the highest column--density component at the edges of the velocity profile, as expected for absorption by a rotating gaseous disk (Wolfe 1995; Prochaska \\& Wolfe 1997). Observations of redshifted 21-cm absorption and emission from DLAs indicate disk--like structures of galactic dimensions (Briggs et al. 1989; Wolfe et al. 1992, Briggs et al. 1997), and Faraday-rotation observations are consistent with the existence of micro-Gauss magnetic fields in these systems (Wolfe, Lanzetta, \\& Oren 1992; Welter, Perry, \\& Kronberg 1984; Perry, Watson, \\& Kronberg 1993). Since the total comoving density of HI in DLAs is comparable to that of stars in the local universe, it is only natural to postulate that the cold gas out of which most of the present population of stars had formed, was already assembled in galaxies at $z\\approx 3$ (Lanzetta et al. 1995; Wolfe et al. 1995; Storrie-Lombardi et al. 1996a, 1996b). However, the typical metallicity of $Z\\approx 10\\% Z_{\\odot}$ (Smith et al. 1996; Lu, Sargent, \\& Barlow 1996, and references therein) and dust-to-gas ratio (Fall \\& Pei 1993; Pei, Fall \\& Bechtold 1991) in these systems imply that star formation was only at its infancy at these early times. This assertion is consistent with recent determinations of the star formation rate at high redshifts (Madau 1996; Lowenthal et al. 1996, and references therein). Since star formation requires cold gas and most of the HI detected through \\lya absorption lines is in damped systems, DLAs are the natural sites for star formation at high redshifts (Fall \\& Pei 1996). In order to unravel the star formation history of the universe, it is of fundamental importance to probe the stars and not only the gas in DLAs. Direct images of distant DLAs typically reveal a luminous galactic core which is separated by $10$--$20~{\\rm kpc}$ from the line-of-sight to the quasar (Steidel et al. 1994, 1995, 1996). However, any inference about the fraction of gas converted into stars in these systems requires a prior knowledge of the Initial Mass Function (IMF) of these stars and their formation history. In addition, a major fraction of the baryonic mass might reside in the outer low surface--brightness halos of these systems. This possibility is raised by recent microlensing searches (Alcock et al. 1996), which indicate that a non--negligible fraction of the massive halo of the Milky--Way galaxy might exist in the form of Massive Compact Halo Objects (MACHOs). As a supplement to local microlensing searches, it would be particularly interesting to examine whether MACHOs populate the halos of galaxies at high redshifts. While the microlensing probability is only $\\sim 10^{-6}$ in the Milky--Way halo ($\\sim 10~{\\rm kpc}$), its value increases up to unity in the cores of halos at cosmological distances ($\\sim 5~{\\rm Gpc}$). This follows from the linear dependence of the lensing cross--section on the observer--lens distance for a source at infinity. Near the center of a DLA, {\\it macrolensing} by the entire galactic potential might take place and yield widely separated quasar images; the likelihood of macrolensing and its generic signatures were discussed recently by Bartelmann \\& Loeb (1996) and by Perna, Loeb, \\& Bartelmann (1997). In this paper we quantify the expected microlensing probability in distant spiral galaxies which show up as DLAs in quasar spectra. The obvious advantage of these galactic systems is that they are selected based on their proximity on the sky to a compact source of light in the form of a quasar. Detection of microlensing of the quasar can be used to probe the stellar {\\it mass} fraction of DLAs and to test whether the halos of galaxies at high redshifts are made of MACHOs. The characteristic Einstein radius of a solar mass lens at a cosmological distance is $\\sim 5\\times 10^{16}~{\\rm cm}$, comfortably in between the scales of the continuum--emitting accretion disk ($\\la 10^{15}~{\\rm cm}$) and the broad line region ($\\sim 3\\times 10^{17}~{\\rm cm}$) of a bright quasar. This implies that a single microlensing event would magnify the continuum but not the broad lines emitted by the quasar. The lines would only be affected by the macrolensing effect of the galaxy as a whole. As a result of this differential amplification, the equivalent width distribution of the broad emission lines (Francis 1992) will be significantly distorted in a sample of microlensed quasars. Microlensing would also result in excess variability of such quasars. In the following, we will quantify the above signatures for a model spiral galaxy which acts as a high redshift DLA. The existence of a population of normal disk galaxies at high redshifts is supported by recent Keck HIRES spectra of the damped \\lya absorber towards the quasar Q 2233+1310 (Lu, Sargent \\& Barlow 1997). The metal absorption lines of the absorbing galaxy are shifted relative to its \\lya emission redshift (Djorgovski et al. 1996), indicating a rotation velocity $\\geq 200~{\\rm km~s^{-1}}$ at a separation $\\sim 20~{\\rm kpc}$. Here we limit our attention to galaxies of this population. For concretness, we focus our discussion on a single spiral galaxy and extrapolate its present--day properties back in time using the simplest closed--box model for its baryonic content. Based on the observed metallicity of DLAs, we assume that the stars constitute only 10\\% of the total disk mass at $z\\approx 3$, and that HI is gradually converted into stars until the present epoch. The galaxy is modeled as having three components: an old stellar bulge, a disk made of gas and stars, and a halo. We compute the significance of microlensing for each of these components. Our simplified model ignores the diversity among different disk galaxies or DLAs and focuses on a typical massive spiral galaxy. Its only purpose is to examine whether microlensing could be detected in this idealized case. Recent imaging of DLA galaxies (Le Brun et al. 1996) indicates that while some systems are clearly spirals, others have irregular morphologies. A more diverse treatment of this problem would involve many more free parameters, and is beyond the scope of this paper. Our primary objective here is to motivate an observational search for these microlensing signatures, rather than to explore the entire possible range of parameter space. The outline of the paper is as follows. In \\S 2 we describe the distributions of HI and stars in our model for a galactic disk. These distributions are then used in \\S 3 to compute the microlensing probability as a function of the intercepted HI column density of the disk. In \\S 4 we repeat the calculation for a halo made of MACHOs. Sections 5 and 6 relate these results to observable signatures. Finally, \\S 7 summarizes our main conclusions. ", + "conclusions": "We have shown that the subset of all quasars which have damped \\lya absorption in their spectra due to an intervening spiral galaxy should possess an enhanched tail in the equivalent width distribution of their broad emission lines, and exhibit excess variability, relative to the rest of the quasar population. Both of these microlensing signatures are more pronounced if the halos of DLAs are composed of compact objects. Previous attempts to find a redshift evolution in the equivalent--width distribution of a large sample of quasars due to microlensing by a population of intergalactic stars have yielded a negative result (Dalcanton et al. 1994). In this paper we have shown that the significance of microlensing relative to the statistical noise should be much more pronounced in a subset of all quasars which are located behind galactic HI disks. Using a simple closed--box model for a spiral galaxy we have found that disk+bulge microlensing could be detected through its imprint on the equivalent--width distortion with a signal-to-noise ratio $\\ga 2$ in a sample of $\\approx 20$ damped \\lya absorbers (DLAs) at $z_{\\rm abs}\\la1$ with $N\\sim 10^{21}{\\rm cm}^{-2}$. The necessary sample size is reduced to $\\approx 10$ if the galaxy halo is composed of MACHOs with a velocity dispersion of $\\sigma\\approx170$ km/sec. The necessary sample size should scale as $\\propto \\sigma^{-4}$ for other galactic systems and therefore massive galaxies are likely to dominate the statistics. In addition, we find that about a tenth of all quasars with DLAs are likely to show excess variability on timescales shorter than five years (cf. Figures 5 and 6). Unfortunately, the current sample of $\\sim 80$ quasars with DLAs (Wolfe et al. 1995) includes only several absorbers with $z\\la1$ and might not be sufficiently large to demonstrate the existence of microlensing. The number of known DLAs could increase by an order of magnitude as a result of spectroscopic follow--ups on the catalog of $\\sim 10^5$ quasars which is expected to be compiled by the forthcoming Sloan Digital Sky Survey (Gunn \\& Knapp 1993). Even before any absorption data is reduced, it would be interesting to select background quasars which are projected close to foreground galaxies and therefore are likely to be microlensed. Related studies by Webster et al. (1988) showed an increased number of quasar-galaxy pairs than expected from random alignments. The enhancement in the quasar surface density near galaxies was interpreted as macro-lensing and used to draw conclusions about the distribution of matter around galaxies. Detection of the level of microlensing for quasars with damped \\lya absorption can be used to calibrate the mass fraction in the form of massive compact objects in galactic halos at high redshifts. Figures 3-6 show that the microlensing signal is enhanced when an isothermal halo made of MACHOs is added to the galactic disk and bulge. Direct imaging of the DLAs can be used to infer the projected separation between the luminous center of the absorber and the quasar. In cases where the DLA redshift is known (e.g. through its \\lya emission; see Djorgovski et al. 1996, and Lu et al. 1997), it might also be possible to infer spectroscopically the velocity dispersion of the intervening galaxy. When combined with the information gathered by microlensing searches in the local universe, such studies could extend our knowledge of the composition of galactic halos out to redshifts as high as $z\\sim 5$." + }, + "9701/astro-ph9701010_arXiv.txt": { + "abstract": "We report the discovery of a probable CH star in the core of the Galactic globular cluster M14 (= NGC 6402 = C1735-032), identified from an integrated-light spectrum of the cluster obtained with the MOS spectrograph on the Canada-France-Hawaii telescope. From a high-resolution echelle spectrum of the same star obtained with the Hydra fiber positioner and bench spectrograph on the WIYN telescope, we measure a radial velocity of $-53.0\\pm1.2$ km s$^{-1}$. Although this velocity is inconsistent with published estimates of the systemic radial velocity of M14 (e.g., ${\\bar {v_r}} \\approx -123$ km s$^{-1}$), we use high-precision Hydra velocities for 20 stars in the central 2\\farcm6 of M14 to calculate improved values for the cluster mean velocity and one-dimensional velocity dispersion: $-59.5\\pm1.9$ km s$^{-1}$ and $8.2\\pm1.4$ km s$^{-1}$, respectively. Both the star's location near the tip of the red giant branch in the cluster color magnitude diagram {\\sl and} its radial velocity therefore argue for membership in M14. Since the intermediate-resolution MOS spectrum shows not only enhanced CH absorption but also strong Swan bands of C$_2$, M14 joins $\\omega$~Cen as the only globular clusters known to contain ``classical\" CH stars. Although evidence for its duplicity must await additional radial velocity measurements, the CH star in M14 is probably, like all {\\sl field} CH stars, a spectroscopic binary with a degenerate (white dwarf) secondary. The candidate and confirmed CH stars in M14 and $\\omega$~Cen, and in a number of Galactic dSph galaxies, may then owe their existence to the long timescales for the shrinking and coalescence of hard binaries in low-concentration environments. ", + "introduction": "Among the presently known types of carbon stars, only the CH stars have abundances and kinematics which are indicative of membership in the Galactic halo (McClure 1985). Since stars of mass M $\\simeq$ 0.8M$_{\\odot}$ are not thought to experience the third dredge-up mechanism during their ascent of the asymptotic giant branch (e.g., Iben 1975), the origin of the enhanced abundances of carbon and {\\sl s}-process elements in these giants remained a puzzle until the discovery that {\\sl all} field CH stars are binaries composed of a red-giant primary and a degenerate (i.e., white dwarf) secondary (McClure 1984; McClure \\& Woodsworth 1990). The peculiar abundances of these objects are therefore thought to be the result of mass exchange via stellar winds or Roche-Lobe overflow during the ascent of the white-dwarf progenitor up the asymptotic giant branch (Han et al. 1995). Clearly, the discovery of {\\sl globular cluster} CH stars would have important implications not only for nucleogenesis and globular cluster abundance anomalies (Pilachowski et al. 1996), but also for the formation, evolution and destruction of binaries in dense environments (Hut et al. 1992). However, searches for CH stars in several globular clusters based on spectroscopic surveys of brightest cluster members (Harding 1962), direct imaging through intermediate-band filters (Palmer 1980; Palmer \\& Wing 1982), and transmission-grating slitless spectroscopy (Bond 1975), have proved, by and large, unsuccessful. At present, a handful of stars having enhanced carbon and {\\sl s}-process elements have been reported in each of $\\omega$~Cen (e.g., Harding 1962; Dickens 1972; Bond 1975), M22 (Hesser, Hartwick \\& McClure 1977; McClure \\& Norris 1977; Hesser \\& Harris 1979), M55 (Smith \\& Norris 1982) and M2 (Zinn 1981). However, while the spectra of these stars {\\sl are} characterized by abnormally high CH absorption compared to other cluster giants, they usually do {\\sl not} show strong Swan bands of C$_2$, suggesting that their anomalous carbon abundances probably arise through a different mechanism, such as incomplete CN processing (Vanture \\& Wallerstein 1992), than that operating in ``classical\" CH stars. Indeed, among this sample of CH-enhanced stars in globular clusters, only two are likely to be genuine CH stars. Both of these stars, RGO 55 (Harding 1962) and RGO 70 (Dickens 1972), are found in $\\omega$~Cen. In this {\\sl Letter}, we report the serendipitous discovery of a carbon star in the core of the poorly-studied globular cluster M14 and argue that it is likely to be a ``classical\" CH star: a post mass-transfer binary consisting of a red giant primary and a white dwarf secondary (McClure \\& Woodsworth 1990). ", + "conclusions": "\\subsection{Membership} From our Hydra spectrum of the carbon star, we measure a radial velocity of $v_r = -53.0\\pm1.2$ km s$^{-1}$ (heliocentric Julian date = 2450282.761). How does this compare to the systemic velocity of M14? Available radial velocity measurements for this cluster have been reviewed by Webbink (1981), who quotes a weighted mean velocity of ${\\bar{v_r}} = -123\\pm5$ km s$^{-1}$. This value is based primarily on low-dispersion spectra of the integrated cluster light obtained by Mayall (1946), as well as a small number of radial velocities (spanning the range +9 to $-$153 km s$^{-1}$) for four Type II Cepheids in the field of M14 accumulated by Joy (1949). An image-tube spectrogram obtained by Hesser, Shawl \\& Meyer (1986) of the integrated cluster light yielded a radial velocity of $v_r = -25\\pm14$ km s$^{-1}$. Given the obvious difficulties with field star contamination in this direction ($l^{\\rm II} = 21.3^{\\circ}$, $b^{\\rm II} = 14.8^{\\circ}$), as well as the large uncertainties of the Mayall (1946) velocities (typical internal error $\\simeq$ 33 km s$^{-1}$, but potentially much larger external errors; see Hesser, Shawl \\& Meyer 1986), we have used our new sample of Hydra velocities to determine improved estimates of the systemic velocity and internal dispersion of M14.\\altaffilmark{10}\\altaffiltext{10}{The result of Hesser, Shawl \\& Meyer (1986) is most likely due to contamination of their spectrum by field stars, since the distribution of radial velocities in this crowded field shows a broad peak near $\\simeq$ $-20$ km s$^{-1}$ (C\\^ot\\'e \\& Welch 1997). Similarly, if the observations of Joy (1949) were made while the Type II Cepheids were near maximum light, then the mean velocity of $-126$ km s$^{-1}$ he obtained may be (partly) explained by atmospheric pulsations, which can approach 30--40 km s$^{-1}$ in stars of this type.} We define a sample of probable M14 members by restricting ourselves to those stars which lie within two half-light radii $r_h$ of the cluster center (i.e., 2$r_h \\simeq$ 2\\farcm6; Trager, King \\& Djorgovski 1995). This sample consists of 20 stars which have a median velocity uncertainty of 1.2 km s$^{-1}$ and span the range $-72.0 \\le v_r \\le -47.6$ km s$^{-1}$. $BV$ photometry for these stars (see Figure 3) demonstrates that they are all likely to be cluster members since they define a smooth red giant branch extending from $V$ = 14.38 to 16.23. From this sample, we find a mean velocity of ${\\bar{{v_r}}} = -59.5\\pm1.9$ km s$^{-1}$ and a one dimensional cluster velocity dispersion of $\\sigma_c = 8.2\\pm1.4$ km s$^{-1}$ using the maximum-likelihood estimators of Pryor \\& Meylan (1993). Based on the close agreement between the radial velocity of the carbon star and that of M14, we conclude that the star is indeed a cluster member. We also note that, if the star in question is similar to the field CH stars studied by McClure \\& Woodsworth (1990), an additional velocity component due to the orbital motion of the red giant primary around the center of mass of the system is expected; among the eight CH stars monitored by McClure \\& Woodsworth (1990), the velocity semi-amplitudes ranged from 4.3 to 12.1 km s$^{-1}$, with a mean of 8.1 km s$^{-1}$. \\subsection{Implications} Radial velocity surveys of field CH stars have provided compelling evidence that {\\sl all} of these systems are post mass-transfer binaries. We therefore conclude that the carbon star in M14 is very likely to be a spectroscopic binary having a white dwarf secondary, although {\\sl confirmation} of duplicity must await additional radial velocity measurements.\\altaffilmark{11}\\altaffiltext{11}{This connection between the field and globular cluster CH stars has been further strengthened by the recent demonstration that both RGO 55 and RGO 70 in $\\omega$~Cen are spectroscopic binaries (Mayor et al. 1996).} It is perhaps notable that, as the only two clusters known to contain candidate or confirmed CH stars, M14 and $\\omega$~Cen are both massive, low-concentration systems. A possible connection between cluster concentration and CH enhancement was noted previously by McClure \\& Norris (1977) who, in a prescient remark, suggested that ``searches for CH stars in the low-concentration clusters M14 and NGC 2419 might be profitable\". Does environment play a role in the evolution of globular cluster CH stars? The dominant dynamical processes affecting binaries in globular clusters are the {\\sl disruption} of ``soft\" binaries through stellar encounters and the {\\sl shrinking} of the orbits of ``hard\" binaries via energy exchanges with intruder stars (e.g., Hut et al. 1992). Equation (1) of Hills (1984) can be used to estimate the separation, $a_s$, of the widest cluster binaries which are expected to have escaped disruption over a Hubble time. Since most Galactic globular clusters have three-dimensional velocity dispersions $\\sigma_v \\lae 15$km s$^{-1}$ (Pryor \\& Meylan 1993), we find $a_s$ $\\gae$ 6 AU. This is comparable to the separation of the {\\sl widest} binaries in the McClure \\& Woodsworth (1990) sample of field CH stars (i.e., 1 $\\lae$ $a$ $\\lae$ 4.5 AU). Therefore, we conclude that the process of CH-star disruption is unlikely to be important for most clusters. On the other hand, the shrinking of hard binaries may play a more significant role. At the cluster center, intruder stars will shrink the orbits of hard binaries at the rate $${{d\\ {\\hbox{ln}}a}\\over{dt}}\\simeq{{-2\\pi Ga{\\rho_c}}\\over{\\sigma_c}}, \\eqno{(1)}$$ where $a$ is the semi-major axis of the binary, ${\\sigma_c}$ is the central one-dimensional velocity dispersion and ${\\rho_c}$ is the central mass density (Hills 1984; Phinney 1996). Integration of this equation yields the initial semimajor axis $a_h$ of the binary whose size is halved over the lifetime, $t_0$, of a cluster: $${a_h} = {{\\sigma_c} \\over 2{\\pi}G{\\rho_c}t_0}. \\eqno{(2)}$$ The orbits of hard binaries that are initially wider than this shrink rapidly, and have $a \\sim a_h$ after a Hubble time. Conversely, tighter binaries evolve on rather longer timescales, and are largely unaffected. Thus, $a_h$ is the characteristic final separation of surviving cluster binaries (Phinney 1996). It is worth bearing in mind, however, that equation (2) refers to the cluster center; the situation for the cluster as a whole is undoubtedly more complicated. In Figure 4 we plot the distribution of $a_h$ against absolute magnitude, $M_V$, for all Galactic globular clusters (filled circles) having measured velocity dispersions and central mass densities (e.g., see Pryor \\& Meylan 1993). Absolute magnitudes are taken from Djorgovski (1993). M14 and $\\omega$~Cen are shown as open stars. The open squares show the location of the Galactic dSph galaxies (Irwin \\& Hatzidimitriou 1995). At least {\\sl some} of the carbon stars seen in several of these galaxies (i.e., Ursa Minor, Draco, Sculptor and Carina) are likely to be CH stars (McClure 1984; Aaronson \\& Olszewski 1987; Armandroff, Olszewski \\& Pryor 1995). The dashed line shows the semi-major axis of the shortest-period binary in the McClure \\& Woodsworth (1990) sample of field CH stars (i.e., $a \\simeq$ 1 AU). This size is similar to the mean radii of low-mass, thermally-pulsating AGB stars (e.g., see Figure 2 of Boothroyd \\& Sackman 1988). Closer initial binary orbits would disrupt the normal evolution of the carbon donor before it reached the thermally-pulsating AGB phase, preventing the dredge-up of carbon. Therefore, Figure 4 suggests that the low concentrations of M14 and $\\omega$~Cen correspond to small values of $\\rho_c/\\sigma_c$, and hence to $a_h$ large enough to accomodate CH systems with separations typical of field CH stars. The discovery of a CH star in M14 therefore provides additional support for the notion (McClure 1984) that the process of CH star formation depends on environment, in particular, on the rate of hard binary shrinking through stellar encounters. Given the incomplete and inhomogenous nature of existing surveys for CH stars in globular clusters, a renewed effort to discover such objects in an expanded sample of clusters (particularly those objects which lie {\\sl above} the dashed line in Figure 4) may provide new insights into the influence of environment on binary evolution." + }, + "9701/astro-ph9701046_arXiv.txt": { + "abstract": "Recent work has shown that the speed of the cooling front in soft X-ray transients may be an important clue in understanding the nature of accretion disk viscosity. In a previous paper (Vishniac and Wheeler 1996) we derived the scaling law for the cooling front speed. Here we derive a similarity solution for the hot inner part of disks undergoing cooling. This solution is exact in the limit of a thin disk, power law opacities, and a minimum hot state column density which is an infinitesimal fraction of the maximum cold state density. For a disk of finite thickness the largest error is in the ratio of the mass flow across the cooling front to the mass flow at small radii. Comparison to the numerical simulations of Cannizzo et al. (1995) indicates that the errors in other parameters do not exceed $(c_{sF}/r_F\\Omega_F)^q$, that is, the ratio of the sound speed at the disk midplane to its orbital velocity, evaluated at the cooling front, to the $q$th power. Here $q\\approx 1/2$. Its precise value is determined by the relevant hot state opacity law and the functional form of the dimensionless viscosity. ", + "introduction": "The most popular model for soft X-ray transients and dwarf novae is that both are due to fluctuations in the luminosity from accretion disks surrounding black holes and white dwarfs, respectively, and that the specific mechanism that starts an outburst is a thermal instability associated with the ionization of hydrogen (for a recent review see Cannizzo 1993a or Osaki 1996). In addition to their intrinsic interest, these systems may provide valuable clues to the nature of angular momentum transport in disks, simply by virtue of being non-stationary systems. In the thermal instability model these disks make the transition back to their quiescent states as their outer parts cool and a thermal transition front propagates inward. Mineshige, Yamasaki, and Ishizaka (1993) pointed out that this model will produce the observed exponential decay of soft X-ray luminosity for soft X-ray transients only if the radius of the hot phase decreases exponentially, i.e. if the cooling front speed is proportional to the radius of the uncooled part of the disk. The speed of the cooling front affects the evolution of the hot phase of the disk and its value will depend on the local physics of the disk, including the dimensionless viscosity $\\alpha$. More specifically, Mineshige (1987) and Cannizzo, Shafter and Wheeler (1988) argued that the cooling front speed is approximately \\begin{equation} \\dot r_F\\approx -\\alpha_F c_{sF}\\left({H\\over r}\\right) \\left({r\\over\\delta r}\\right), \\label{eq:first} \\end{equation} where $\\alpha_F$ is the dimensionless viscosity at the onset of rapid cooling, $H$ is the disk thickness, $\\delta r$ is the width of the cooling region, $c_s$ is the sound speed in the midplane of the disk, and the subscripts `F' denote values at the onset of rapid cooling. Numerical simulations seemed to show that $\\delta r\\sim 0.1 r$. Recent high resolution work by Cannizzo, Chen and Livio (1995, hereafter CCL) has shown that this estimate for $\\delta r$ was a consequence of insufficient resolution in the simulations. They found $\\delta r$ was approximately equal to the geometric mean of the disk height and radius. Combining this with equation (\\ref{eq:first}) they argued that an exponential decline implied a scaling law for $\\alpha$ of the form \\begin{equation} \\alpha\\approx 50\\left({c_{sF}\\over r_F\\Omega_F}\\right)^{3/2}\\propto \\left({H\\over r}\\right)^{3/2}. \\end{equation} More recently Cannizzo (1996) has shown that a similar result can be recovered from models of the dwarf nova SS Cygni. In a previous paper (\\cite{VW96}, hereafter VW) we showed that the cooling front speed was actually independent of the structure of the disk outside the radius where rapid cooling sets in, and that the cooling front speed scales as \\begin{equation} \\dot r_F\\propto-\\alpha_F c_{sF}\\left({H\\over r}\\right)^q, \\label{eq:second} \\end{equation} where $q$ is determined by the functional form of $\\alpha$ and the opacity law for the hot part of the disk. Assuming that the hot phase is thermally stable implies that $q$ lies in the interval $[0,1]$. For most reasonable opacity laws $q\\approx 1/2$. This implies that the exponential luminosity decline of dwarf novae and soft X-ray transients is not just another test of the form of $\\alpha$, but one that specifically depends only on the physics of the hot state, thereby eliminating uncertainties regarding opacities in the cold state and the effects of comparing a largely ionized hot state to a largely neutral cold one. Finally, in VW we managed to show that the cooling front speed is surprisingly insensitive to the value of $\\alpha$ right at the cooling front, but instead is a measure of its behavior over a broad annulus inward from the cooling front. While it is important to see that the scaling law reported in CCL can be recovered from the simple physical argument presented in VW. It would be better still to see that the numerical coefficients found by CCL can be reproduced by an analytic solution built around this same argument. Furthermore, the numerical solution is not precisely self-similar and by itself the scaling law does not give us an understanding of when it might fail. In this paper we present a self-similar analytic solution of a cooling wave in a thin accretion disk and estimate the error involved in applying it to a disk of finite thickness. Section 2 contains a derivation of the thin disk similarity solution. Section 3 contains a summary of our results, a discussion of the error involved in applying this solution to realistic disks, and a discussion of the implications of this work for the physics of accretion disks. ", + "conclusions": "We have found a similarity solution that describes the progress of a cooling wave in an infinitely thin accretion disk whose hot state is governed by power law opacities. For a disk of finite thickness, the nonzero value of the critical column density for the onset of rapid cooling introduces an error at the outer edge of the similarity solution. The free parameters of this solution are almost entirely determined by the physics of accretion and the hot state opacity, the sole exception being the parameter $\\Delta$, which is determined by the structure of the rapid cooling zone. In practice, this means that at present $\\Delta$ is determined from numerical simulations. We can see, by comparison with the numerical simulations of CCL, that if we take a Kramers opacity and $n=3/2$ we get a solution which is a fair description of the progress of a cooling wave in a disk surrounding a soft X-ray transient. Looking at our results in figure 1, we can see that the temperature profile of the hot phase is almost a perfect power law, with a sharp cutoff near the outer edge. Clearly, it will be difficult to observe the cutoff structure directly. However, this model predicts that the power law temperature profile of the inner region will be well above the critical temperature for the thermal instability at the cooling front. This should be a testable prediction of this model, although this effect will be more dramatic in the later stages of the luminosity decline. Conversely, a significantly broader temperature transition profile than the one seen here would be an indication of nonlocal effects in $\\alpha$. In order to evaluate the size of the errors induced by using this solution, we need to examine the structure of the cooling wave near the cooling front. In this region the state variables will have a radial scale length much less than $r$, so the dynamical equations can be approximated by saving only radial derivatives of $\\dot M$ and $\\Sigma$. With this in mind we can write \\begin{equation} \\partial_t\\Sigma(r)\\approx -\\dot r_F\\partial_r\\Sigma. \\end{equation} Invoking equation (\\ref{eq:contin}) we have \\begin{equation} -\\dot r_F\\partial_r\\Sigma={-1\\over 2\\pi r }\\partial_r\\dot M. \\end{equation} We can integrate this from $\\Sigma=0$ to $\\Sigma=\\Sigma_F$ to obtain \\begin{equation} \\Delta\\dot M\\approx -\\dot r_F 2\\pi r_F\\Sigma_F={-\\dot r_F\\over V_F}\\dot M_F, \\label{eq:sigext} \\end{equation} where $\\Delta \\dot M$ is the amount by which the mass outflow at the edge of the similarity solution, where $\\Sigma=0$, exceeds the mass outflow from the actual cooling front, where $\\Sigma=\\Sigma_F$. {}For the particular case examined by CCL we have \\begin{equation} \\Delta\\dot M\\approx 5.4 \\left({c_{sF}\\over r_F\\Omega_F}\\right)^q\\dot M_F, \\end{equation} which will be somewhat less than half for $(c_{sF}/r_F\\Omega_F)$ between $10^{-2}$ and $10^{-3}$. Although this amounts to a large correction to $\\dot M$, in agreement with the comparison between CCL and the prediction of the similarity solution, it does not imply that the behavior of $\\dot M$ with time is subject to a similarly large correction. Assuming a Kramers law opacity and $n=1.5$ we can look at the evolution of the mass flow through the point where $\\Sigma=\\Sigma_F$. We find that \\begin{eqnarray} t\\partial_t\\ln(\\dot M_F-\\Delta\\dot M)&=t\\partial_t \\ln\\left[\\dot M_0 \\left(1-{0.94\\over\\Delta}\\left({c_{sF}\\over r_F\\Omega_F}\\right)^q\\right)\\right]\\cr &=t\\partial_t\\dot M_0-{0.94\\over\\Delta}t\\partial_t\\left({c_{sF}\\over r\\Omega_F}\\right)^q\\cr &=t\\partial_t \\ln\\dot M_0\\left[1-{1.88C_1\\over\\Delta}r_F\\partial_{r_F} \\left({c_{sF}\\over r_F\\Omega_F}\\right)^q\\right]\\cr &=t\\partial_t\\ln\\dot M_0\\left(1-0.6\\left({c_{sF}\\over r_F\\Omega_F}\\right)^q\\right). \\end{eqnarray} In other words, we get only a small error when we treat the time evolution of the mass flow through the $\\Sigma=\\Sigma_F$ radius as though it were strictly proportional to the mass flow across the outer surface of the similarity solution. Since the realistic condition on the mass flux is that it is fixed at $\\Sigma=\\Sigma_F$, rather than at $\\Sigma=0$, the negative sign in the above equation actually implies that the realistic solution will give $ t\\partial_t \\ln \\dot M_0$ slightly more than the one predicted from the similarity solution. For the calculation of CCL the discrepancy should be of order a few percent, which is not significant. Similarly we can estimate the error we make in equating the radius of the cooling front with the $\\Sigma=0$ surface by integrating equation (\\ref{eq:torque}) using equations (\\ref{eq:mdotdef}) and (\\ref{eq:codef}). We obtain \\begin{equation} \\Delta r = r_F {3\\over\\Delta} {c_{sF}\\over r_F\\Omega_F}. \\end{equation} {}For the calculation of CCL this implies an error of a few percent. CCL pointed out that comparing the observed e-folding time for the soft X-ray luminosity in X-ray transients to the results of their simulation gives a dimensionless viscosity of \\begin{equation} \\alpha\\approx 50 \\left({h\\over r}\\right)^{3/2}. \\end{equation} This coefficient of $50$ is surprisingly large, since one would expect that almost any theory which gave the correct scaling law would involve a coefficient of order unity. This estimate of the coefficient is roughly inversely proportional to the minimum hot state temperature, which is fixed by the physics of the opacity in the hot state. The only other obvious way to get a different result is by using a different value of $\\Delta$, the ratio of the outflow velocity at the cooling front to $\\alpha_F c_{sF}$. However, we see from equation (\\ref{eq:rfdot}) that this enters into the expression for the speed of the cooling front only as $\\sim\\Delta^{1/2}$. Even if we replace the value calculated by CCL, $\\Delta\\sim 1/6$, with $\\Delta=1$ the estimate of the coefficient in the expression for $\\alpha$ only drops to $\\sim 20$. One last point we need to consider is under what circumstances we can ignore the structure of the cold region in discussing the velocity and structure of the cooling front. From the equation of continuity we see that the column density in the cold disk region, just outside the cooling front, will be of order $\\sim \\Sigma_F (V_F/-\\dot r_F)$. If this exceeds the maximum stable column density for the cold gas, then the solution described here is inapplicable, and the cold region will affect the velocity and structure of the cooling wave. The value of the maximum cold state column density is somewhat uncertain. Here we will use the results of Cannizzo (1993b). Assuming $n\\approx 3/2$, we get \\begin{equation} {\\Sigma_{c,max}\\over \\Sigma_F}=6.2 \\left({\\alpha_F\\over 0.1}\\right)^{-0.31} r_{10}^{-0.036}, \\end{equation} where $r_{10}$ is the front radius in units of $10^{10}$ cm. We need this to be larger than $V_F/-\\dot r_F$ which is \\begin{equation} {V_F\\over -\\dot r_F}\\approx \\Delta \\left({50\\over\\alpha_F}\\right)^{q/n}. \\end{equation} We see from these expressions that this condition is generally satisfied, although not by a large factor, with very little dependence on $\\alpha_F$ or $r_{10}$. Whether or not it will be satisfied for substantially different models for $\\alpha$ or disk opacities depends on the specifics of the models. In sum, the similarity solution described in this paper will give a good description of the progress of a cooling wave in a hot disk as long as four conditions are met. First, the outflow velocity across the cooling front needs to be substantially greater than the cooling front velocity. In practice, this means that the hot part of the disk needs to have a height to radius ratio of order $10^{-2}$ or less for the thermal instability associated with hydrogen ionization. If this condition is violated then the errors in the solution can become large. Other thermal instabilities will involve slightly different conditions, depending on the values of $q$ and $\\Delta$. Second, the the opacity of the hot disk material, for a few e-foldings interior to the cooling front, has to be described by a power law. The explicit solution here assumed a Kramers law, but any power law approximation will yield a similar solution. Third, the $\\alpha$ needs to be purely a function of local disk parameters. This last condition is significant, inasmuch as the only prediction for $\\alpha$ which self-consistently allows for the effects of local magnetic instabilities and gives the correct scaling law is based on an explicitly non-local model (cf. \\cite{VJD90}, \\cite{VD92}). We plan an extension of this work to cover this model and to see if an observably different cooling front structure can be recovered from it. Fourth, neglecting the structure of the disk outside the radius where rapid cooling sets in is appropriate only when the ratio of $V_F$ to $\\dot r_F$ is not so large as to imply an unphysically large column density in the cold gas. This does not seem to be an important limitation for the systems considered to date, but this conclusion is sensitive to the details of the cold disk structure." + }, + "9701/astro-ph9701052_arXiv.txt": { + "abstract": "We simulate the dynamical and chemical evolution of a dwarf galaxy embedded in a dark matter halo, using a three-dimensional $N$-body/SPH simulation code combined with stellar population synthesis. The initial condition is adopted in accord with a $10^{10}M_\\odot$ virialized sphere in a $1\\sigma$ CDM perturbation which contains 10\\% baryonic mass. A supersonic spherical outflow is driven by the first star burst near the center of the galaxy and produces an expanding super shell in which stars are subsequently formed. Consecutive formation of stars in the expanding shell makes the stellar system settled with the exponential brightness profile, the positive metallicity gradient, and the inverse color gradient in agreement with observed features of dwarf galaxies. We therefore propose that the energy feedback via stellar winds and supernovae is a decisive mechanism for formation of less compact, small systems like dwarf galaxies. ", + "introduction": "Our understanding of how galaxies originate and distribute on large scales in the universe has greatly improved in the last two decades. While the standard model of hierarchical galaxy clustering \\markcite{WR1987}(White \\& Rees 1978) has been successful in explaining the clustering pattern of galaxies revealed by redshift surveys, it predicts a large number of low-mass galaxies ($L<10^{10}L_{\\odot}$) beyond that estimated from the observed luminosity function of galaxies (\\markcite{WF1991}White \\& Frenk 1991; \\markcite{CAFNZ1994} Cole {\\it et al.} 1994). The hierarchical model should therefore involve some mechanism which suppresses the formation of such small galaxies. Main mechanisms so far proposed include an energy feedback from supernovae that prevents the collapse of a forming galaxy (Dekel \\& Silk 1986; Lacey \\& Silk 1992) and a photoionization by ultraviolet background radiation that keeps the gas hot and unable to collapse (Dekel \\& Rees 1987; Efstathiou 1992). Significant body of new observations of nearby dwarf galaxies has reveald a web of filaments, loops and expanding super giant shells which are imprinted in the ionized gas around individual galaxies (\\markcite{MFD1992}Meurer, Freeman \\& Dopita 1992; \\markcite{MHW1995}Marlowe, Heckman \\& Wyse 1995; \\markcite{H1996}Hunter 1996). Since the traces of energetic winds are oriented from supernovae or massive stars, it is evident that the heat input from them greatly affects the dynamics of small galaxies. This feedback of energy into the interstellar medium must play a decisive role in the early stage of galaxy evolution when star formation rate is expected to be much higher. Dekel \\& Silk (1986) showed that the supernova feedback mechanism nicely accounts for the observed correlations between metallicity, color and luminosity of galaxies (see also Vader 1986; Yoshii \\& Arimoto 1987). There is however a clear distinction in structural and chemical quantities bewteen dwarf ellipticals (dEs) and normal ellipticals in spite of their morphological similarity. The central concentration of dEs is relatively low and their luminosity profiles are best fitted by an exponential function, whereas the profiles of normal ellipticals are known to follow de Vaucouleurs' law (\\markcite{FL1983}Faber \\& Lin 1983; \\markcite{BST1984}Binggeli, Sandage \\& Tarenghi 1984; \\markcite{IWO1986}Ichikawa, Wakamatsu \\& Okamura 1986; \\markcite{CB1987}Caldwell \\& Bothun 1987). Moreover, the color of many dEs becomes redder towards outer radii of the system (Vader {\\it et al.} 1988; Kormendy \\& Djorgovski 1989; Chaboyer 1994), and this trend of color gradient is clearly opposite to normal galaxies. The origin of these striking features of dEs remains yet to be explained (for a review see Ferguson \\& Binggeli 1994). In particular, no attempts have ever been made to examine whether the supernova feedback mechanism is viable also in this context. In this paper, we use three dimensional simulation code with a cosmologically motivated initial condition and investigate the formation and evolution of a dE galaxy taking into account the dynamical responses of the system from supernova-driven winds. Our simulation shows that such winds propagating outwards from inside the system collide with the infalling gas and produce the super shell in which stars are formed. This specific process of star formation turns out to reproduce the observed features of dEs, and therefore the heating by supernovae proves to be an ideal suppressing mechanism against the efficient formation of low-mass galaxies in the hierachical clustering model. ", + "conclusions": "A three-dimensional $N$-body/SPH simulation code, combined with stellar population synthesis, is used to follow the dynamical and chemical evolution of a dwarf protogalaxy with $10^{10}M_\\odot$ (baryonic/dark=1/9) which originates from a $1\\sigma$ CDM perturbation. This less massive galaxy receives significant dynamical responses from the heat input by stellar winds and supernovae. The first star burst near the center of the system produces a supersonic spherical outflow of the gas. This outflow collides with the infalling gas and gives rise to an expanding dense shell. Then, stars begin to form in the expanding shell with its site propegating outwards with the shell. We find from the simulation that this consecutive process of star formation creates the exponential brightness profile and the inverse color gradient of the system in agreement with the observations of dwarf galaxies. \\markcite{A1994}Athanassoula (1994) performed one-dimensional simulations of the dynamical evolution of dE galaxies including the energy feedback from supernovae. The models without dark halo are shown to give a better agreement with observations than those with dark halo. Our more realistic, three-dimensional simulations however indicate that the dark halo is necessary and plays a vital role to form the bound stellar system, otherwise the system is blown out to disrupt completely. In general, the color gradient of galaxies is created by the gradient in either metallicity or age of the underlying stellar population. Simple models of chemical evolution of galaxies usually predict the negative metallicity gradient which corresponds to the color becoming redder towards the galaxy center. Since dwarf galaxies have the inverse color gradient, Vader {\\it et al.} (1988) were led to interpret this observed trend in terms of the positive age gradient. We note however that stars with very low metallicities must have been formed on very short timescales and therefore no appreciable age difference results. The above puzzling situation indicates that previous results based on simple models of chemical evolution can not be applied to small systems like dwarf galaxies. We demonstrate in this paper that dynamical modelling is the only proper way to investigate the evolution of dawrf galaxies. Successful reproduction of their basic features in our simulation suggests that the stellar energy feedback mechanism is indeed a likely mechanism against the efficient formation of low-mass galaxies in the CDM universe." + }, + "9701/astro-ph9701114_arXiv.txt": { + "abstract": "We study the early evolution of the electron fraction (or, alternatively, the neutron-to-proton ratio) in the region above the hot proto-neutron star formed after a supernova explosion. We study the way in which the electron fraction in this environment is set by a competition between lepton (electron, positron, neutrino, and antineutrino) capture processes on free neutrons and protons and nuclei. Our calculations take explicit account of the effect of nuclear composition changes, such as formation of alpha particles (the \\lq\\lq alpha effect\\rq\\rq ) and the shifting of nuclear abundances in nuclear statistical equilibrium associated with cooling in near-adiabatic outflow. We take detailed account of the process of weak interaction freeze-out in conjunction with these nuclear composition changes. Our detailed treatment shows that the alpha effect can cause significant increases in the electron fraction, while neutrino and antineutrino capture on heavy nuclei tends to have a buffering effect on this quantity. We also examine the effect on weak rates and the electron fraction of fluctuations in time in the neutrino and antineutrino energy spectra arising from hydrodynamic waves. Our analysis is guided by the Mayle \\& Wilson supernova code numerical results for the neutrino energy spectra and density and velocity profiles. ", + "introduction": "\\label{sec:introye} In this paper we examine the early evolution of the electron fraction, $Y_e$, in the post-core-bounce supernova environment. The electron fraction is defined to be the net number of electrons per baryon: \\begin{equation} \\label{eq:defye} Y_e = (n_{e^-} - n_{e^+}) / n_b = 1 /(1 + N_n/ N_p); \\end{equation} where $n_{e^-}$, $n_{e^+}$, and $n_b$ are the proper number densities of electrons, positrons, and baryons, respectively. The latter expression in equation (\\ref{eq:defye}) follows from the condition of overall charge neutrality. Here, $N_n$ and $N_p$ are the total proper neutron and proton densities, respectively, so that $n_b = N_n + N_p$ and $N_n / N_p$ is the net neutron-to-proton ratio. The electron fraction is a crucial determinant of nucleosynthesis produced from neutrino-heated ejecta in models of post core-collapse supernovae (cf., \\markcite{Woosley92} Woosley \\& Baron 1992; \\markcite{Meyer92} Meyer et al. 1992; \\markcite{Woosley94} Woosley et al. 1994; \\markcite{Woosley92} Woosley \\& Hoffman 1992; \\markcite{Qian93} Qian et al. 1993; \\markcite{Qian96} Qian \\& Woosley 1996; \\markcite{Wiit93} Witti, Janka \\& Takahashi 1993; \\markcite{Hoffman96a} Hoffman et al. 1996a; \\markcite{Hoffman96b} Hoffman et al. 1996b). We follow the terminology of \\markcite{Fuller92} Fuller, Mayle, Meyer, \\& Wilson (1992), \\markcite{Fuller93} Fuller (1993) and \\markcite{Qian95} Qian \\& Fuller (1995) and divide the post core bounce evolution of the outflowing material above the nascent neutron star into two epochs: (1) the shock reheating or \\lq\\lq $p$-Process\\rq\\rq\\ epoch at times post-core-bounce $t_{pb} \\lesssim \\, 1 \\, {\\rm s};$ and (2) the neutrino-driven wind or \\lq\\lq $r$-Process\\rq\\rq\\ epoch at $t_{pb} \\gtrsim \\, 1 \\, {\\rm s}.$ We expect the shock reheating epoch to be characterized by chaotic outflow (\\markcite{Burrows95}Burrows, Hayes, \\& Fryxell, 1995; \\markcite{Miller93} Miller, Wilson \\& Mayle 1993; \\markcite{Janka95}Janka \\& M\\\"uller 1995; Herant, Benz, Hix, Fryer \\& Colgate 1994; Janka \\& M\\\"uller(1996)) and rapid heating by neutrino interactions of material behind the shock. Neutrino-heated ejecta originating from this epoch have been suggested to give rise to the neutron number N=50 peak $r$-Process material (\\markcite{woosley92}Woosley \\& Hoffman 1992; \\markcite{Meyer92}Meyer et al. 1992; \\markcite{woosley94}Woosley et al. 1994) and possibly at least some of the light $p$-Process nuclei such as $^{92}{\\rm Mo}$ and $^{96}{\\rm Ru}$ (\\markcite{Fuller95}Fuller \\& Meyer 1995; \\markcite{Hoffman96a} Hoffman et al. 1996a). However, in the \\markcite{Woosley94}Woosley et al. (1994) calculations (based on the Mayle \\& Wilson supernova results) the N=50 $r$-Process nuclides originating in this epoch are grossly overproduced relative to solar system abundances. Two fixes have been proposed for the \\lq\\lq N=50 overproduction problem\\rq\\rq: (1) Fuller \\& Meyer (1995) invoke a modification of linear rapid outflow, a high neutrino fluence, and the alpha effect to increase $Y_e$ and thereby reduce $N=50$ overproduction; and (2) Hoffman et al. (1996a) show that as long as the electron fraction at this epoch can be engineered to be $Y_e \\gtrsim 0.484$, N=50 overproduction is avoided. As a bonus, both of these fixes concomitantly suggest that some of the light $p$-nuclei will be synthesized. Hoffman et al. (1996a) find that the light $p$-nuclei are produced in the correct proportions so long as $0.484 \\lesssim Y_e \\lesssim 0.488$. This epoch is characterized by relatively low entropy per baryon, $s/k \\sim 40$, and relatively higher $Y_e$ compared to the conditions which may obtain at $t_{pb} \\gtrsim 1 \\, {\\rm s}$. Note that the Hoffman et al. (1996a) work implies that we may have to compute $Y_e$ to of order $1\\%$ accuracy to predict confidently the nucleosynthesis. It could be that this will not be necessary, as the hydrodynamic outflow is phased in just the right way that a given mass element experiences the necessary $Y_e$ regime at the necessary temperature. Only future computations can address this issue. As we will show, for a given outflow model, predicting $Y_e$ histories to $1\\%$ accuracy may be next to impossible at this stage given our crude understanding of neutrino transport and multidimensional hydrodynamic effects (cf. Herant, Benz, Hix, Fryer \\& Colgate 1994; Janka \\& M\\\"uller 1996). By contrast, the later $r$-Process epoch is characterized by considerably higher entropy, $s/k \\approx 100-500$ (see \\markcite{Qian} Qian \\& Woosley 1996 and Meyer et al. 1992), and possibly by a well ordered, near steady state, outflow resembling a neutrino-driven wind (\\markcite{Duncan86}Duncan, Shapiro, \\& Wasserman 1986, \\markcite{Meyer92} Meyer et al. 1992, \\markcite{Qian96} Qian \\& Woosley 1996). In fact, \\markcite{Woosley94} Woosley et al. (1994) have shown that the bulk of the solar systems' $r$-Process material with nuclear mass $A \\gtrsim 100$ could be synthesized in this epoch. However, considerable controversy surrounds the theoretical modeling of conditions in the \\lq\\lq hot bubble\\rq\\rq\\ which forms in this regime. Though the \\markcite{Woosley94} Woosley et al. (1994) calculations yield a near perfect solar abundance distribution for the heavier $r$-Process nuclides, they are based on the very high entropy ($s/k \\sim 400$) conditions obtained by the Wilson and Mayle supernova code. Not only have such high entropies been challenged (\\markcite{Qian96}Qian \\& Woosley 1996 find $s/k \\lesssim 200$), but even if the entropy were $s/k \\gtrsim 300$, neutrino neutral current spallation of alpha particles (\\markcite{Meyer95}Meyer 1995) has been shown to result in a drastic and fatal reduction in the neutron-to-seed ratio required for the $r$-Process. Though models show the material in the hot bubble to be quite neutron rich, $Y_e \\approx 0.4$, \\markcite{Hoffman1996b} Hoffman et al. (1996b) and \\markcite{Luo} Meyer, Brown \\& Luo (1996) have demonstrated that far lower values of $Y_e$ are necessary to obtain the requisite neutron-to-seed ratio for the $r$-Process if the entropy is $s/k \\lesssim 200$. At entropies this low, the bad effects of neutrino spallation of alpha particles \\markcite{Meyer95} (Meyer 1995) would be evaded. Hoffman et al. (1996b) discuss the neutron-to-seed ratio and $Y_e$ issues related to this epoch, while \\markcite{Fuller96} Fuller, Qian, \\& Wilson (1996) and \\markcite{Caldwell96} Caldwell, Fuller \\& Qian (1996) discuss neutrino flavor mixing schemes which could give lower $Y_e$ and hence help the $r$-Process. In this paper we perform a complimentary study of the evolution of the electron fraction which concentrates on the effects of nuclear composition changes. We focus in particular on the early time, \\lq\\lq low\\rq\\rq\\ entropy environment of the shock reheating epoch. In what follows we concentrate on the weak interaction balance essentially in a single outflowing mass element (i. e., one-dimensional outflow). We employ outflow results from the Mayle \\& Wilson supernova code to illustrate various effects bearing on $Y_e$. It should be kept in mind that we are not {\\it predicting} $Y_e$, as the Mayle and Wilson results may not be representative of the true picture of supernova evolution. For example, multidimensional hydrodynamic effects could effectively cause different mass elements to have different time/thermodynamic histories. Nevertheless, we choose the simplest case (1D outflow) to elucidate the physics. In section \\ref{sec:over}, we present an overview of all of the variables which affect the calculation of the electron fraction. We show explicitly how the various charged current lepton capture processes are important. In section \\ref{sec:nucap}, we discuss salient aspects of the electron neutrino and antineutrino capture rates on free nucleons. We explore the difference between using the Mayle and Wilson transport calculation-derived neutrino energy spectra as an example and approximate blackbody spectra. We also discuss the effect on the electron fraction of hydrodynamic wave-induced fluctuations in the neutrino energy spectra. In section \\ref{sec:elcap}, we assess the role of electron and positron capture processes on $Y_e$. In section \\ref{sec:alpha}, we examine the \\lq\\lq alpha effect\\rq\\rq\\ or the tendency of the formation of alpha particles to raise the electron fraction. We discuss the equilibrium and nonequilibrium nature of weak interaction freeze-out in section \\ref{sec:noneq}, with particular attention to the role of nuclear composition changes and the role of neutrino capture on heavy nuclei. We give conclusions in section \\ref{sec:conclusion}. ", + "conclusions": "\\label{sec:conclusion} In this paper we have given an in depth treatment and analysis of the evolution of the electron fraction in neutrino-heated supernova outflow, including detailed treatment of the effects of nuclear composition changes. Our study has concentrated on the time $t_{pb} \\lesssim 1 \\, {\\rm s}.$ The evolution of $Y_e$ is this epoch may be quite important for models of the light $p$-Process and the neutron number $N \\approx 50$ $r$-Process nuclei (Hoffman et al. 1996a; Fuller \\& Meyer 1995). We have employed fits to the detailed neutrino and antineutrino energy spectra from the Mayle and Wilson supernova calculations. We find that these detailed spectra are necessary for computations of weak rates to the level of accuracy in $Y_e$ which may be required for nucleosynthesis considerations. However, we find that hydrodynamic wave-induced fluctuations in the ratios of neutrino and antineutrino spectral parameters with time produce significant excursions in $Y_e$. We find that the rates of electron and positron capture on free nucleons can also be important for computing the evolution of the electron fraction. The charged current weak rates freeze out from equilibrium at a time when the electron and positron capture rates may still have some influence on $Y_e$. During this time period, the effect of these rates is to increase the electron fraction. We have given detailed calculations of the \\lq\\lq alpha effect\\rq\\rq\\ - the increase in the electron fraction caused by a changing alpha particle mass fraction. Our results indicate that the alpha effect can play a very significant role in setting $Y_e$. We have employed numerical calculations of nuclear composition changes in nuclear statistical equilibrium, coupled with fits to density and velocity of outflow histories from the Mayle and Wilson results, to compute the evolution of $Y_e$. These calculations explicitly include all charged current weak interaction processes, including neutrino and antineutrino capture on heavy nuclei. We follow the evolution of $Y_e$ through the epoch of weak equilibrium freeze out. These calculations show that the combination of neutrino and antineutrino capture on heavy nuclei and antineutrino capture on free protons tend to keep $Y_e$ constant. The results presented in this paper are meant to illustrate several different variables and processes which can alter the electron fraction in post core bounce supernova outflow. Clearly, more sophisticated models of neutrino transport and hydrodynamic outflow than those employed here would be necessary to actually {\\it predict} $Y_e$ in a reliable fashion. We believe, however, that the effects described here will always play the major role in setting $Y_e$." + }, + "9701/astro-ph9701028_arXiv.txt": { + "abstract": "{ Taking advantage of the advances in array detector technology, an imaging polarimeter (IMPOL) has been constructed for measuring linear polarization in the wavelength band from 400-800~nm. It makes use of a Wollaston prism as the analyser to measure simultaneously the two orthogonal polarization components that define a Stoke's parameter. An achromatic half-wave plate is used to rotate the plane of polarization with respect to the axis of the analyser so that the second Stoke's parameter also can be determined. With a field of view correponding to about $\\rm 30\\times30~mm^2$ for a \\diameter 1.2~m, f/13 telescope, a sensitive, liquid-$\\rm N_2$ cooled CCD camera as the detector and a built-in acquisition and guidance unit, the instrument can be used for studying stellar fields or extended objects with an angular resolution of $\\sim $2\\arcsec\\ . The instrumental polarization is less than 0.05\\% and the accuracies of measurement are primarily limited by photon noise for typical observations.} ", + "introduction": "The advances in two dimensional array detector technology in the optical and near infrared wavelength bands have made new kinds of imaging astronomical observations feasible. Astronomical polarimetry is one field which has gained tremendously from these developments. The limitations of using aperture photometry for polarimetry were so severe that any serious study was rendered time consuming and difficult. On the other hand, imaging polarimetry with its capabilities for multiplexing, simultaneous sky measurement, seeing-limited resolution etc. offer great advantages over aperture polarimetry. Astronomers have recognized this potential and have developed several new observation techniques in the optical (eg. Scarrott 1991; Jannuzi et. al. 1993; Jarrett et. al. 1994; Wolstencroft et. al. 1995; Simmons~et.~al. 1995) and near infrared (eg. Kastner \\& Weintraub 1994; Moore \\& Yamashita 1995; Weintraub et. al. 1995) wavelenghts to study phenomena in a variety of Galactic and extragalactic astrophysical objects. In this paper, we report the design and construction of an imaging polarimeter (IMPOL) which uses a cooled CCD array as detector. It was developed at the Inter-University Centre for Astronomy and Astrophysics \\hbox{(IUCAA)}, INDIA. The principle of the instrument (Sen \\& Tandon 1994) is based on a combination of ideas suggested by Ohman (1939) and Appenzeller (1967). An instrument of this type has been constructed at the University of Durham and has been in use for some time now (Scarrott~et. al. 1983). Section~2 is a description of the instrument -- the different subsections dealing with various aspects ranging from design guidelines to instrument control and user-interface. The dominant sources of errors in the measurement are investigated in Sect.~3, while an estimate of the performance of the instrument under two typical observing conditions is given in Sect.~4; Section~4 also contains an estimate of the performance of the acquisition \\& guidance unit. In Sect.~5 we discuss the observational procedure and the data-analysis method. Section~6 contains results of the commissioning tests of the instrument. The last section (Sect. 7) contains the concluding remarks. ", + "conclusions": "An Imaging Polarimeter (IMPOL) has been constructed which uses mostly standard optical and electrical components, but through careful design has been able to achieve \\medskip \\pscaption{\\psboxto(\\hsize;0cm){fig5.ps}}{Fig. 5 Normalized Stoke's parameter $q$ is plotted against $u$ for a number of unpolarized standard stars. The correlation coefficient of the points is about 0.06 and the average values of $q$ and $u$ give a value of $p = 0.03\\%$.} \\medskip \\noindent photon-noise limited performance. The instrument has a field of view of about 6.5\\arcmin\\ for a \\diameter 1.2~m, f/13 telescope, and the detector is a sensitive, liquid N$_2$ cooled CCD. An off-axis acquisition and guidance unit capable of using stars as faint as $V_{\\rm mag}\\sim 15$ is also built in the instrument so that long exposures of faint extended objects like reflection nebulae etc. can be taken while keeping the image fixed on the CCD face to within one half of a pixel -- this stability of the image is a necessary condition to achieve a high accuracy in relative photometry between the different frames used to estimate the polarization. Observations of nearby standard polarized and unpolarized stars show that for wideband observations, there is no discernible depolarization and the instrumental polarization is less than 0.05\\%. Preliminary results of wideband polarimetry of stellar fields, using a PSF fitting technique for determining the centroids of the stellar images and aperture photometry with a radius of 3 pixels to derive the intensities of the individual images, give close to photon-noise limited accuracy of 0.15\\% for $\\rm V_{mag} = 15$ stars with about an hour of total exposure time. For extended objects with brightness about 20th mag. per square arcseconds and a background of about the same brightness, it is estimated that an accuracy of about 1\\% is possible with 60 minutes of total exposure time and an aperture of 10~sq. arcsec." + }, + "9701/astro-ph9701245_arXiv.txt": { + "abstract": "We have observed the $^{12}$CO J=2-1, $^{13}$CO J=2-1, and $^{12}$CO J=3-2 lines in a sample of seven giant molecular clouds in the Local Group spiral galaxy M33 using the James Clerk Maxwell Telescope. The $^{12}$CO/$^{13}$CO J=2-1 line ratio is constant across the entire sample, while the observed $^{12}$CO J=3-2/J=2-1 line ratio has a weak dependence on the star formation environment of the cloud, with large changes in the line ratio seen only for clouds in the immediate vicinity of an extremely luminous HII region. A large velocity gradient analysis indicates that clouds without HII regions have temperatures of 10-20 K, clouds with HII regions have temperatures of 15-100 K, and the cloud in the giant HII region has a temperature of at least 100 K. Interestingly, the giant HII region appears capable of raising the kinetic temperature of the molecular gas only for clouds that are quite nearby ($< 100$ pc). The continuous change of physical conditions across the observed range of star formation environments suggests that the unusual physical conditions in the cloud in the giant HII region are due to post-star formation changes in the molecular gas, rather than intrinsic properties of the gas related to the formation of the giant HII region. The results from this study of M33 suggest that similar observations of ensembles of giant molecular clouds in more distant normal spiral galaxies are likely to give meaningful measurements of the average physical conditions inside the molecular clouds. These results also imply that clouds with a factor of three difference in metallicity have similar density and temperature, which in turn imply that the differences in the CO-to-H$_2$ conversion factor seen in these clouds can be attributed entirely to metallicity effects. ", + "introduction": "Determining the physical conditions inside molecular clouds is important for understanding the link between the properties of the molecular gas and the types and amount of stars that are formed. Cloud properties that could affect the star formation process include the temperature and density of the molecular gas, as well as the mass fraction in high density gas. For example, higher gas temperatures might be required to form high-mass stars (\\markcite{t84}Turner 1984), while a cloud with a higher mass fraction of dense gas might form stars with a higher efficiency (e.g. \\markcite{L91}Lada et al. 1991). In return, star formation, particularly massive star formation, can affect conditions inside molecular clouds by compressing the gas at the boundaries of stellar wind or supernova shocks and heating the gas by increasing the ultraviolet radiation field. Giant HII regions are particularly interesting targets because they concentrate many hundreds of OB stars in a relatively small volume of space and represent the only nearby prototypes for the extreme star formation conditions in starburst and ultraluminous galaxies. In our own Galaxy, observations of low-luminosity HII regions suggest that single high-mass stars can form from relatively low-mass molecular clouds (\\markcite{hm90}Hunter \\& Massey 1990). Observations of several nearby giant molecular clouds show significant differences in the spatial distribution of star formation. In the Taurus-Auriga region, star formation occurs throughout the cloud as single stars or in small clusters (\\markcite{k90}Kenyon et al. 1990), while Orion B has most of its star formation occurring in four massive dense cores (\\markcite{L91}Lada et al. 1991), and Orion A and Ophiuchus are each forming most of their stars in a single dense cluster (\\markcite{ms94}McCaughrean \\& Stauffer 1994, \\markcite{wly89}Wilking, Lada, \\& Young 1989). Despite the differences in the spatial distribution of the young stars, the star formation efficiency is 1-4\\% for all four regions (\\markcite{el91}Evans \\& Lada 1991). Interestingly, the star formation efficiency in two giant HII regions in M33 is also in the range of 2-5\\% (\\markcite{wm95}Wilson \\& Matthews 1995). The clumpy structure of molecular clouds (e.g. \\markcite{SG90}Stutzki \\& G\\\"usten 1990) implies that the volume-averaged density of a cloud is not a good measure of the density of the molecular gas producing most of the observed emission. For extragalactic molecular clouds, their low and usually unknown filling factor within the beam also implies that the kinetic temperature cannot be estimated from the peak brightness temperature of the observed line. Instead, the density and temperature must be determined using radiative transfer models and observations of many line ratios of CO and its isotopomers. Because this technique requires observations of some of the rarer isotopomers such as $^{13}$CO, it was first applied to starburst galaxies with strong CO lines (e.g. \\markcite{t91}Tilanus et al. 1991, \\markcite{G93}G\\\"usten et al. 1993). However, one difficulty with even the nearest starburst galaxies is that the observations measure the emission from many molecular clouds within a single beam. For example, at a distance of 3 Mpc, a 20\\arcsec~ beam subtends a diameter of 290 pc, which is much larger than the 10-100 pc diameter of giant molecular clouds in the Milky Way (\\markcite{sss85}Sanders, Scoville, \\& Solomon 1985). If the physical conditions in molecular clouds vary significantly from one cloud to the next, the results of fitting the average line emission may bear little resemblance to even the average properties of the molecular cloud ensemble. Local Group galaxies offer an advantage for such studies because they are close enough (50 kpc - 1 Mpc) that it is possible to isolate {\\it individual} molecular clouds within a 15-20\\arcsec~ beam. In addition, individual clouds can cover a substantial fraction of the beam and thus the CO lines are reasonably strong. Observations of a large sample of individual clouds can identify variations in the density and temperature from one cloud to the next, as well as correlations in cloud properties with the presence and intensity of massive star formation. The uniformity (or lack thereof) of the molecular cloud population also provides a means to test the reliability of radiative transfer techniques for determining density and temperature in more distant galaxies, where only the average emission from the molecular cloud population can be observed. In addition, density and temperature measurements for individual clouds can be used to assess whether changes in the density and brightness temperature of the gas are likely to produce systematic errors in the observed correlation of the CO-to-H$_2$ conversion factor with metallicity (\\markcite{w95}Wilson 1995, \\markcite{dss86}Dickman, Snell, \\& Schloerb 1986, \\markcite{s96}Sakamoto 1996). In this paper we present observations of seven giant molecular clouds in the spiral galaxy M33. The clouds were chosen to cover a wide variety of star formation conditions, from clouds with no optical HII regions to a cloud located in the brightest giant HII region in the galaxy. The clouds also cover a range of three in oxygen abundance (\\markcite{v88}Vilchez et al. 1988). The observations and data reduction are discussed in \\S 2. The observed line ratios are compared with previous Galactic and extragalactic observations and with the star formation environment of the clouds in \\S 3. The density and temperature obtained from an analysis of the line ratios with a large velocity gradient code are presented in \\S 4. The implications of the results for observations of more distant galaxies and the calibration of the CO-to-H$_2$ conversion factor as a function of metallicity are discussed in \\S 5. The paper is summarized in \\S 6. ", + "conclusions": "We have observed the $^{12}$CO J=2-1, $^{13}$CO J=2-1, and $^{12}$CO J=3-2 lines in a sample of seven giant molecular clouds in the Local Group spiral galaxy M33. The clouds were chosen to cover a range of star formation conditions, from clouds without optical HII regions to a cloud in a giant HII region. We find that the $^{12}$CO/$^{13}$CO J=2-1 line ratio is constant across the entire sample, while the $^{12}$CO J=3-2/J=2-1 line ratio is somewhat smaller in three clouds without optical HII regions than in three clouds with HII regions. The seventh cloud, located in the brightest giant HII region in the galaxy, has an even higher $^{12}$CO J=3-2/J=2-1 line ratio. We conclude that the $^{12}$CO J=3-2/J=2-1 line ratio of a cloud has a weak dependence on the star formation environment of the cloud, with large changes in the line ratio seen only for clouds in the immediate vicinity of an extremely luminous HII region. We used a large velocity gradient code to determine the density, temperature, and column density for the clouds. The analysis indicates that clouds without HII regions have temperatures in the range of 10 to 20 K and clouds with HII regions have temperatures in the range of 15 to 100 K, while the cloud in the giant HII region has a kinetic temperature of at least 100 K. We note that the giant HII region seems to have a relatively limited sphere of influence within which it can heat the molecular gas, since a molecular cloud located only 120 pc from the giant HII region shows a normal cooler temperature. The column density increases by about an order of magnitude from the cool clouds to the hot cloud, as does the volume filling factor of the dense gas, while the density decreases by less than an order of magnitude. The continuity of physical properties (kinetic temperature, density, column density, and filling factor) across the range of star formation environments in our sample suggests that the unusual physical conditions seen in the cloud in the giant HII region are due to post-star formation changes in the molecular gas, rather than intrinsic properties of the gas related to the formation of the giant HII region. Excluding the giant HII region, the relatively uniform line ratios observed in M33 suggest that average physical conditions determined from similar measurements of ensembles of giant molecular clouds in more distant normal spiral galaxies are likely to give physically meaningful results. These uniform line ratios also imply that the average brightness temperature and density are similar in clouds ranging over a factor of three in metallicity, which in turn suggests that differences in the CO-to-H$_2$ conversion factor seen in these clouds can be attributed entirely to metallicity effects." + }, + "9701/astro-ph9701135_arXiv.txt": { + "abstract": "A study of the ultraviolet continuum variability (Paltani \\& Courvoisier \\cite{PC94}) has shown that the relative variability of quasars and Seyfert galaxies decreases when the luminosity increases. The spectral information included in the IUE spectra allows us to study this dependence in the rest frame of the objects. The trend is strengthened by the general property that active galactic nuclei vary more at short wavelengths than at long wavelengths in the ultraviolet domain. The scatter observed in all other studies is still present. An important part of this scatter may however be explained if one tries to estimate the uncertainties on the variability due to the sampling. We discuss the variability using the concept of discrete events. The trend between variability and luminosity is described by a power-law with an index $-0.08$, which is incompatible with the power-law of index $-1/2$ predicted by the most general discrete-event models in which the change in average luminosity is due to differences in average event rates exclusively. Several biases are investigated, but we conclude that the $-1/2$ index is definitely inconsistent with the data. A flat relationship is however possible, if some bias has been underestimated. We propose different ways whereby discrete events may produce a different variability--luminosity relationship: by changing the luminosity or the life time of the events, or by introducing interdependence between the events. The latter possibility cannot produce a satisfactory relationship. Using the former possibilities, we do not find any ``natural'' explanation for the variability--luminosity relationship in the context of discrete-event models. This is possibly an indication that explanations in which variability is not expressed in terms of discrete events should be favoured. ", + "introduction": "The statistical properties of the ultraviolet variability of active galactic nuclei (AGN) have been investigated by Paltani \\& Courvoisier (\\cite{PC94}) (hereafter PC94). This study, which was based on spectra obtained by the {\\em International Ultraviolet Explorer} (IUE) from the ULDA ({\\em Uniform Low Dispersion Archive}) database, has been done in the observer's frame. In addition to the variability at 2\\,000 \\AA, we quantified also the change of variability with the wavelength, which showed that all classes of AGN have a larger variability at short wavelength than at long wavelength in the IUE domain. In this paper we shall consider only the objects that have been called \\SF\\ in PC94, i.e.\\ Seyfert 1 galaxies, radio-quiet quasars, and low-polarization radio-loud quasars. The relationship between variability and luminosity in quasars and Seyfert galaxies has been investigated by many authors. Most of them found an anti-correlation between these two quantities (Pica \\& Smith \\cite{PS83}; Cristiani et al.\\ \\cite{Cal90}; PC94; Hook et al.\\ \\cite{Hal94}), while others found a correlation or no correlation at all (Bonoli et al.\\ \\cite{Bal79}; Tr\\`evese et al.\\ \\cite{Tal89}; Giallongo et al.\\ \\cite{Gal91}). Rediscussion of these last three papers by Hook et al.\\ (\\cite{Hal94}) seems to show that all the samples are compatible with the existence of an anti-correlation. All these studies have used estimates of the variability in the observer's frame, as they can be much more easily obtained than those in the objects' rest frame. Therefore the correct relationship between variability and luminosity is still unknown. This relationship is very important, as it can be compared with the predictions of some AGN models, in particular the starburst model of Terlevich (\\cite{T92}). In this paper we establish this relationship in the objects' rest frame, and we discuss some biases that can affect it. We then check the compatibility between our result and the predictions of these models. We formalize in the most general way the concept of ``discrete-event models'', i.e.\\ models where the total light curve is the sum of a series of light curves, each one associated with one ``event'', whatever its physical nature, and we compare their predictions with the above constraint. We give the general relationship between variability and luminosity for these models. We study two different ways by which this relationship can be modified. In the first one we assume that some of the event parameters may actually depend on the luminosity. In the second one we examine the effects of the introduction of interdependence between the events. ", + "conclusions": "The trend between the variability and the luminosity seen by many authors is confirmed when one estimates the rest-frame variability. The most probable index of the relationship is $-0.08$. It is compatible with results obtained using high-luminosity quasars. This is an additional argument for the continuity between the Seyfert galaxies and the most powerful quasars. We have investigated the effects of several biases, and we consider our result as very probably correct. An intrinsic $\\eta=-1/2$ index, which is what is generally expected from discrete-event models, would be possible only if there exists very large and yet undiscovered biases; but this seems highly unlikely. The other extreme, which is that the variability does not depend on the luminosity, is more plausible. However the difficulty in producing such a relationship with discrete events is not significantly alleviated. The easiest way to produce a relationship with an index different from $-1/2$ is to assume that the average event luminosity increases with luminosity. In this case, it would mean that the event luminosity must increase only about 15 \\% more slowly than the total luminosity. The scaling of the event stretching factor may also contribute to the relationship. The relationships given here may provide useful constraints for any kind of discrete-event model. The possibility that the birth time of the events does not follow a Poisson distribution has also been investigated and cannot be completely excluded. However the system that we have explored cannot (alone) produce the correct relationship. Anyway a non-Poisson system appears somehow unnatural, and the physical situation at its origin might be difficult to find. The predictions of physical models that can be expressed in terms of discrete events can be compared quantitatively with the variability--luminosity relationship using Eq.~(\\ref{eq-an_var_ind}). But, it appears to us that discrete-event models do not satisfy the variability--luminosity relationship in a natural way. This difficulty possibly means that this kind of model does not apply to the active galactic nuclei, and, therefore, that Eq.~(\\ref{eq-dis_ev}) is not valid. Other source of variability (not considering the possibility that variations are extrinsic, which seems very unlikely) could be global effects, which do not have a simple general analytical formulation, in opposition to the local processes discussed here." + }, + "9701/astro-ph9701192_arXiv.txt": { + "abstract": "We present the results of two complementary ground-based programmes to determine the host galaxy properties of radio-quiet and radio-loud quasars and to compare them with those of radio galaxies. Both infrared images and optical off-nuclear spectra were obtained and we discuss the various strategies used to separate the quasar-related emission from that of the underlying galaxy. However, the key feature of this project is the use of carefully matched samples, which ensure that the data for different types of object are directly comparable. ", + "introduction": "The paper briefly describes the results of a continuing long-term project to study the host galaxies of powerful AGN. The aim of the project is two-fold: to test the scheme for radio-loud quasars (RLQs) and radio galaxies (RGs) which attemps to unify the two types of object via orientation effects, and to investigate the extent to which the host galaxy influences the radio properties of the AGN by comparing the hosts of radio-loud and radio-quiet quasars (RQQs). Using ground--based observations we have approached the question of host galaxy properties from two independent directions: near-infrared ($K$-band) imaging, to determine the host morphologies and luminosities, and off-nuclear optical spectroscopy to investigate their star-formation histories. A more detailed description of the near-infrared imaging can be found in Dunlop {\\it et al.} (1993) and Taylor {\\it et al.} (1996). ", + "conclusions": "This has proved to be a very fruitful project. Many interesting (and some unexpected) results have emerged from the near-infrared imaging study and, although the off-nuclear spectroscopy is still very much a work in progress, the fact that we have been able to isolate starlight in all of the spectra taken so far is a very encouraging result. A consistent picture is emerging from the data. It appears that RLQ hosts and RGs are indeed the same type of galaxy - large luminous spheroidal systems with old, red stellar populations - consistent with the unified scheme. The hosts of RQQs are also large and luminous, and can be either disc-dominated or spheroidal systems. There seems to be a tendency for the most luminous RQQs to occur in elliptical hosts. Ages of the RQQ hosts cover a wide range and the bluest galaxies all seem to be undergoing interactions. In the immediate future we have been awarded 34 orbits on the HST to observe our three AGN samples in $R$-band. Not only will these images enable us to determine the optical morpholgies of the hosts with a level of detail which is impossible from the ground, but by providing us with reliable optical luminosities for the host galaxies they will enable us to bridge the gap between our two ground-based datasets, allowing us to calculate $R-K$ colours for the galaxies and thus to test whether a particular spectrophotometric model can explain the shape of a galaxy spectrum from optical through to near-infrared wavelengths." + }, + "9701/astro-ph9701121_arXiv.txt": { + "abstract": "Results of the X-ray spectral analysis of the high-redshift radio-quiet quasar Q1101-264 are presented. The ASCA spectrum suggests a marginal evidence of a FeK$\\alpha$ emission line at about 2 keV (observer's frame). Both the ASCA and ROSAT spectra are well fitted by a power law with spectral slope $\\Gamma$$\\sim$1.9. This is the first 0.3-30 keV spectrum (rest frame) of a z $>$ 2 radio-quiet quasar. ", + "introduction": "Q1101-264 is a high-redshift (z=2.15) radio-quiet quasar. It was observed with the gas imaging spectrometer (GIS) and solid state spectrometer (SIS) on board the ASCA satellite (Tanaka et al. 1994) in June 1996 and with the Rosat PSPC (Pfeffermann et al. 1987) in December 1993. SIS grade 6 data were also included in order to improve the statistics above 5-6 keV (Mukai \\& Weaver 1996). The GIS data had too low statistics for a spectral analysis and were, therefore, excluded from the following analysis. The total exposure times after screening were $\\sim$ 17 Ks/SIS and 5 Ks for ROSAT PSPC. ", + "conclusions": "" + }, + "9701/astro-ph9701067_arXiv.txt": { + "abstract": "The operation of CGRO/BATSE continues to produce, after more than 5 years, a valuable database for the study of long-term variability in bright hard X-ray sources. The all-sky capability of BATSE provides, using the Earth occultation technique, up to $\\approx$30 flux measurements per day for each source. The long BATSE baseline and the numerous rising and setting occultation flux measurements allow searches for periodic and quasi-periodic signals from hours to hundreds of days. We present initial results from our study of the hard X-ray variability in 24 of the brightest BATSE sources. Power density spectra are computed for each source. In addition, we present profiles of the hard X-ray orbital modulations in 8 X-ray binaries (Cen X-3, \\mbox{Cyg X-1,} Cyg X-3, GX 301-2, Her X-1, OAO1657-415, Vela X-1 and 4U1700-37), several-hundred-day modulations in the amplitude and width of the main high state in the 35-day cycle in Her X-1, and variations in outburst durations and intensities in the recurrent X-ray transients. Keywords: X-ray binaries; Long-term monitoring; CGRO/BATSE ", + "introduction": "The continued operation of the BATSE experiment on CGRO provides a valuable resource for the study of variability in hard X-ray sources. The BATSE experiment's Large Area Detectors (LADs) consist of eight separate NaI(Tl) scintillation detectors positioned in an octahedral pattern to provide continuous coverage over 4$\\pi$ steradians (see Fishman et al. 1989). The detectors are sensitive from 20-1800 keV. Although uncollimated, the large collecting area, 2025 $\\rm{cm^2}$ each, and the low-Earth orbit of CGRO provide the opportunity to monitor various discrete X-ray objects using occultations by the Earth (see Harmon et al. 1992). Continuous monitoring also allows the study of very bright and pulsed sources on time scales shorter than the CGRO orbit (91-94 min) and the 8 Spectroscopy Detectors are available for coverage down to around 8 keV, but we limit this analysis to LAD occultation measurements of bright sources. The ability of BATSE to provide essential information to \\mbox{INTEGRAL} on variable sources is discussed elsewhere in these proceedings (see Fishman et al. 1997). ", + "conclusions": "" + } +} \ No newline at end of file