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Using the same telescope settings, we also obtained the spectra of the AGB stars IRC+10216 (catalog ) and CRL 3068 (catalog ) , the PPN CRL 2688 (catalog ) , and the young PN NGC 7207 (catalog ) . A detailed study of the chemical compositions in different evolutionary stages will be published in a separate paper. |
We are grateful to the ARO staff for their help during the observing run. Otto Peng and Yu-Chin Huang assisted in the observations. We also thank Jun-ichi Nakashima and Jin-Hua He for useful discussions. We acknowledge an anonymous referee for comments which helped strengthen the paper. The work was supported by a gran... |
# Source: arxiv 0808.3282 # Title: Most Population III Supernovae are Duds # Sections: all # Downloaded: 2026-03-02T07:58:02.138919+00:00 |
MOST POPULATION III SUPERNOVAE ARE DUDS Robert L. Kurucz Harvard-Smithsonian Center for Astrophysics August 22, 2008 MOST POPULATION III SUPERNOVAE ARE DUDS |
Robert L. Kurucz Harvard-Smithsonian Center for Astrophysics 60 Garden St, Cambridge, MA 02138 Abstract One Population III dud supernova produces enough oxygen to enable 10 solar masses of primordial gas to bind into M dwarfs. This is possible because radiation from other Population III stars implodes the mixture of ox... |
Subject headings: supernovae — dark matter Introduction Most of the physics and literature on Population III star formation and on Population III supernovae are speculative, as is this paper. The goal of supernova modelling is always to produce supernova explosions and products. Methods and physics that do not work to ... |
Population III stars, dud supernovae, supernovae, oxygen dwarfs In prestellar times small perturbations cool by radiating in HD, LiH, and other light molecules and start to collapse. If the perturbations are too massive or have too much angular momentum, they cannot successfully collapse to form a star. |
At some threshold of mass or angular momentum, a small perturbation [MATH] 1000 M ) can collapse by somehow ejecting enough mass and angular momentum to produce a Population III star of, say, 200 M |
that is rapidly rotating and oblate. As it evolves, the star continues to lose a significant fraction of its mass and angular momentum through a radiatively-driven wind from the loosely bound mass at the equator. The star produces a huge prolate “Stromgren sphere”. The core of the star is rapidly rotating, oblate, and ... |
over its lifetime that is only on the order of 10 years. It burns to an oxygen core, but then, because of the rapid rotation, the oblateness, and the rapid evolution, the star cannot stably collapse. It flies apart in the attempt. It explodes with, say, 10 -3 of the energy of a supernova and produces ejecta that move a... |
Larger perturbations produce (several) smaller Population III stars with less angular momentum. One of these stars, with 150 to 100 M evolves more slowly and stably than a more massive star. It is less rapidly rotating and less oblate. It loses a significant fraction of its mass and angular momentum through a radiative... |
When radiation from a Population III star hits an oxygen blob from a dud supernova, it compresses the gas and the gas radiates in OH and H O to cool rapidly. Ultraviolet radiation from the star is down-converted to visible-, infrared-, and radio-line radiation that cannot be absorbed by the primordial gas surrounding t... |
I have computed ATLAS12 (Kurucz 1995; 2005) model atmospheres for oxygen dwarfs with Teff 3500K and 3000K; log g = 5; for fractional abundances by number H = .911, He = .089, log Li = -10, log O = -4,-5,-6. The pinch of Li adds a few electrons. Figure 1 shows the temperature-Rosseland optical depth relations compared t... |
Let there be, say, 10 dud supernovae per large protogalaxy. The protogalaxies become full of oxygen blobs. Let there be, say, 10 Population III stars smaller than 150 M that radiate 10 51 ergs/M and then supernova. If the Population III stars and the oxygen blobs are randomly positioned relative to each other in the pr... |
As the oxygen blobs are used up by forming stars, they are replaced by the Population III supernova remnants that have iron or alpha elements or both. Supernova remnants are energetic and mix through much larger volumes of primordial gas than the oxygen blobs were able to; the average abundances are much lower. The adm... |
The galaxy-size perturbations that fill the universe are themselves filled with oxygen-dwarf and Population II globular clusters, with oxygen-dwarf and Population II field stars, and with black holes from Population III supernovae. Most of the gas has been formed into stars. The evolutionary details are discussed in my... |
Prediction The “dark matter” halo consists of 10 , or so, oxygen-dwarf globular clusters, up to 25 per square degree. The clusters would appear 10” to 60” in diameter and would have thousands of stars per square arcsecond, but would still be transparent. The observational test is to produce [M RED – M IR ] color maps i... |
References Kurucz, R.L. 1995. A new opacity-sampling model atmosphere program for arbitrary abundances. in Stellar Surface Structure, ed. K.G. Strassmeier and |
J.L. Linsky (Dordrecht: Kluwer), 523-526. Kurucz, R.L. 2000. An outline of radiatively driven cosmology. (astro-ph/0003381). Kurucz, R.L. 2005. Memorie della Societa Astronomica Italiana Supplementi 8, |
10-20. Figure captions Figure 1. Temperature-log Rosseland optical depth relations for model atmo- spheres with Teff 3500K, log g 5.0 . The colors are the same in all the figures. Red is Population I solar abundance. Green is Population II 1/300 solar abundance with alpha elements enhanced by 2.5 . Cyan is an oxygen dw... |
Figure 2. Energy distributions at resolving power 500 for the wavelength range 0 to 3000 nm. The smooth curves represent the continuum level. The models shown are Population I, Population II, and oxygen dwarf -4. Full scale is 1.5 [MATH] 10 ergs cm -2 -1 nm -1 at the star. |
Figure 3. As Figure 2 but for the wavelength range 3000 to 10000 nm. Full scale is 7.5 [MATH] 10 ergs cm -2 -1 nm -1 at the star. |
Figure 4. As Figure 2 but for oxygen dwarfs -4, -5, and -6. Figure 5. As Figure 4 but for the wavelength range 3000 to 10000 nm. Figure 6. As Figure 4 but for Teff 3000K, log g 5.0 . Figure 7. As Figure 6 but for the wavelength range 3000 to 10000 nm. |
# Source: arxiv 0808.3299 # Title: The Dark-Matter Fraction in the Elliptical Galaxy Lensing the Quasar PG1115+080 # Sections: all # Downloaded: 2026-03-02T07:58:03.245997+00:00 |
The Dark-Matter Fraction in the Elliptical Galaxy Lensing the Quasar PG 1115+080 Abstract We determine the most likely dark-matter fraction in the elliptical galaxy quadruply lensing the quasar PG 1115+080 based on analyses of the X-ray fluxes of the individual images in 2000 and 2008. Between the two epochs, the [MATH... |
1. Introduction The theory of gravitational lensing is by now quite well understood (e.g., the review by Narayan & Bartelmann, 1999 . For the case of a quasar quadruply imaged by an intervening galaxy, a very simple model for the lensing potentials — a monopole plus a quadrupole — usually succeeds in fitting the positi... |
There are two leading explanations for this small scale structure. One intriguing explanation is that we are seeing milli -lensing by dark matter condensations of sub-galactic mass (Witt et al., 1995 ; Mao & Schneider, 1998 ; Dalal & Kochanek, 2002 ; Metcalf & Madau, 2001 ; Chiba, 2002 , which are predicted in large nu... |
If the flux ratio anomalies are due to milli -lensing, i.e., [MATH] [MATH] [MATH] dark matter condensations (Wambsganss & Paczyński, 1992 , then the Einstein radii of such masses, projected back to the quasar, are sufficiently large that we would expect the flux ratios to (i) be the same at all wavelengths, and (ii) re... |
In fact, neither of these expectations based on milli -lensing is observed, and the results are overwhelmingly more compatible with stellar micro -lensing. If this is indeed the case, then (i) the stellar Einstein radii projected back to the quasar are more nearly comparable in size with the expected quasar emission re... |
Recently we have systematically analyzed ten quadruply-imaged quasars using Chandra X-ray Observatory archival data and Hubble Space Telescope visible images (Pooley et al., 2007 . We find that the flux ratio anomalies in the X-ray images of quads are systematically larger than for the same quads imaged in the visible,... |
(Pooley et al., 2007 . Thus, the observationally inferred optical emission regions in quasars, based on microlensing, are much larger than anticipated. This is an intriguing mystery to be pursued. |
Of equal astrophysical significance, the same observations of the amplitudes and frequency of occurrence of X-ray flux ratio anomalies can also be used to infer the fraction of dark matter at distances from the center of the lensing elliptical galaxies corresponding to the impact parameter of the images (typically [MAT... |
2. Observations of PG 1115+080 2.1. Prior Observations of PG 1115+080 PG 1115+080 was the second gravitationally lensed quasar to be discovered (Weymann et al., 1980 and the first one found to be quadruple. It has been the subject of numerous studies at wavelengths ranging from radio to mid-infrared to optical to UV to... |
From its discovery more than a quarter century ago to the present, the optical flux ratio between images [MATH] and [MATH] has been in the range of [MATH] 0.65–0.85 as determined from numerous measurements (see Fig. ), and Vanderriest et al. ( 1986 reported [MATH] varying by [MATH] 30% on a timescale of one year via me... |
2.2. 2008 Chandra observation of PG 1115+080 PG 1115+080 was observed for 28.8 ks on 2008 January 31 (ObsID 7757) with the Advanced CCD Imaging Spectrometer (ACIS). The data were taken in timed-exposure mode with an integration time of 3.24 s per frame, and the telescope aim point was on the back-side illuminated S3 ch... |
Reduction was performed using the CIAO 4.0 software provided by the Chandra X-ray Center. The data were reprocessed using the CALDB 3.4.3 set of calibration files (gain maps, quantum efficiency, quantum efficiency uniformity, effective area) including a new bad pixel list made with the acis_run_hotpix tool. The reproce... |
Our analysis follows the procedure laid out in Pooley et al. ( 2007 . We produced a 0.3–8 keV image of PG 1115+080 with a resolution of 0.0246 ′′ per pixel. To determine the intensities of each lensed quasar image, a two-dimensional model consisting of four Gaussian components plus a constant background was fit to the ... |
In order to visualize the dramatic rise in the flux of [MATH] (with the [MATH] and [MATH] images clearly separated), we produced maximum likelihood reconstructions of two Chandra images from the 2000 and 2008 observations. For this, we used the max_likelihood function in the IDL Astronomy User’s Library, which is based... |
The long-term history of the [MATH] flux ratio in the optical and X-ray is summarized in Fig. . The X-ray history of the fluxes from each of the quasar images, based on our two-dimensional Gaussian fits and the measured spectrum of image [MATH] , is given in Table and displayed in Fig. , showing that the change in [MAT... |
3. Evaluation of Dark-to-Stellar Matter Content The way in which observations of flux ratio anomalies, and their variations with time, can lead to an estimate of the dark-to-stellar matter content of the lensing elliptical galaxy is based on analyses of stellar microlensing magnification maps (Wambsganss, 1999 . The ma... |
We approach the magnification map analyses in three slightly different and complementary ways, described and discussed below. 3.1. Bayesian analysis 1 |
The distributions of magnifications produced from such maps, for two different values of dark-to-stellar matter ratios, are shown in the top panels of Fig. (the histograms have been shifted, which we describe below). The different colored histograms in each panel are for the HS, HM, LS, and LM images. These histograms ... |
We treat the histograms in Fig. as likelihood distributions of the intrinsic source intensity, given the observed fluxes of the four individual images (during the 2000 observation). Without loss of generality, we take the zeropoint to be the intrinsic source intensity such that the observed flux for image [MATH] is exa... |
We have repeated the same calculations for nine different stellar fractions. The results are shown in Fig. as relative likelihood plotted against stellar fraction. Note that the scale of stellar fraction is not linear. From this result we conclude that the most likely fraction of mass in stars in the lensing galaxy at ... |
3.2. Bayesian analysis 2 We take a slightly different approach here to mitigating our ignorance of the intrinsic intensity of the source. We perform an analysis based on each image’s fraction of the total observed intensity (which is based on our two-dimensional fitting described in §2.2). The total model intensity is ... |
We use Bayes’s theorem to calculate [EQUATION] where [MATH] is the observed flux (expressed as a fraction of the total intensity) of image [MATH] |
[MATH] ) in the 2000 observation and [MATH] is the stellar fraction for which a particular microlensing map (and its associated fractional intensity map) is generated. We discuss each of the three terms on the right-hand side. |
The term [MATH] represents the probability to obtain the observed intensity in 2000 for a specific stellar fraction. To calculate this, we count all pixels in the fractional intensity map for that stellar fraction which lie within the 1- [MATH] confidence interval of [MATH] (e.g., [MATH] ) and divide by the total numbe... |
The term [MATH] is the prior probability on a specific stellar fraction, which we take to be uniform. All of these values are therefore [MATH] |
The denominator, [MATH] , is similar to the first term but without regard to any particular stellar fraction. We calculate this by counting the number of pixels in all maps that lie within the 1- [MATH] confidence interval of [MATH] and dividing by the total number of pixels in all maps. |
Using this framework, we calculate the posterior probability of each stellar fraction based on each value of [MATH] . We then multiply those probabilities for each stellar fraction together to produce a plot similar to Fig. . These results are shown in Fig. . The probabilities based on the individual images (left panel... |
Because the previous analysis and this one utilize similar methods, one might expect them to yield identical answers. Indeed, the right panel of Fig. shows good agreement with the first Bayesian approach, namely, the most likely breakdown of matter in the lensing galaxy at the typical impact parameter of the four image... |
3.3. Bayesian analysis 3 The third approach utilizes two epochs of Chandra data. We calculate the posterior probability of each stellar fraction given the observations in 2000 and 2008, subject to the constraint that the source and magnification map moved a certain amount relative to each other in the intervening eight... |
We use Bayes’s theorem to calculate [EQUATION] We again discuss the three terms on the right-hand side. The term [MATH] represents the probability to obtain the observed intensity in 2008 for a specific stellar fraction and observed intensity in 2000. To calculate this, we first find all pixels, for a given stellar fra... |
The term [MATH] is the left-hand side of eq. ( ) and was explained in the previous section. The denominator, [MATH] , is similar to the first term but without regard to any particular stellar fraction. We simply sum over all maps and calculate this term as [MATH] |
Similar to the previous analysis, we calculate the posterior probability of each stellar fraction based on each [MATH] . We then multiply those probabilities for each stellar fraction together to produce a plot similar to Fig. , and this is shown in Fig. . Comparing this with the previous Bayesian analysis of just the ... |
A more comprehensive approach to multiply sampled lightcurves has been developed by Kochanek ( 2004 who then applied it to the case of Q 2237+0305. The locations of the images of this quasar are close to the core of the (barred spiral) lensing galaxy because the lens redshift is low. This makes it much more likely that... |
3.4. Evaluation of rapid temporal change We use the stellar microlensing magnification maps to investigate the likelihood that the [MATH] image would have changed its intensity by a factor of [MATH] 6 during an interval of eight years (in the observer frame). We choose this interval because it is the length of time bet... |
3.5. Discussion of the Bayesian methods The individual panels of the left-hand side of Fig. are essentially another way of looking at the magnification map histograms in the top panels of Fig. ; the panels of Fig. show the probabilities of the nine magnification maps to produce the flux fraction that was observed. Wher... |
When the additional information that [MATH] became much less demagnified on a timescale of eight years is added, the combined probability becomes less strongly peaked. The individual panels on the left-hand side of Fig. show that, in comparison to the 2000 data alone, this difference is due mainly to the difference in ... |
If the relative motion between the source and the magnification map were much larger, it would be easier to accommodate both the 2000 and 2008 observations of [MATH] with the low-stellar-fraction maps, as it would be if the temporal baseline were much larger, which it very well could be. Fig. 3 of Pooley et al. ( 2006 ... |
In addition to the uncertainties in baselines (due to sparsely sampled light curves) and relative velocities, there is another shortcoming to this type of Bayesian analysis, namely, the use of static magnification maps to analyze temporal behavior. As shown by Kundic & Wambsganss ( 1993 ); Wambsganss & Kundic ( 1995 , ... |
Given the additional complications introduced by the motions of the quasar and the microlensing stars, the first and second Bayesian analyses are clearly more straightforward, free from the uncertainties in the third. In a forthcoming paper, we will apply the first method to 14 quadruply lensed quasars that Chandra has... |
Finally, we point out one common feature of all three approaches: the important role that the low-magnification images ( [MATH] and [MATH] ) play in constraining the stellar fraction. This is evident in the individual panels of Figs. and . It can also be seen in the upper panels of Fig. by considering the effects of th... |
4. Summary and Conclusions We have observed a dramatic change in the X-ray flux of the [MATH] image of PG 1115+080 in Chandra observations that were separated by eight years. The short timescale for the flux change clearly indicates the presence of stellar microlensing rather than milli-lensing due to dark matter subst... |
One particularly interesting aspect of the observed change in flux ratio between the [MATH] and [MATH] images is the very rapid timescale on which it occurred ( [MATH] 8 years). Figure indicates how far bulk translation of the lensing galaxy is likely to move the quasar image with respect to the caustic pattern during ... |
D. P. thanks Nicholas E. Matsakis and Eric M. Downes for extremely useful discussions of Bayesian analysis and is grateful to the MIT Kavli Institute for Astrophysics and Space Research for its hospitality during the summer of 2008 during which most of this work was done. D. P. and S. R. gratefully acknowledge support ... |
# Source: arxiv 0808.3334 # Title: Chromospheric magnetic field and density structure measurements using hard X-rays in a flaring coronal loop # Sections: all # Downloaded: 2026-03-02T07:58:04.565739+00:00 |
Chromospheric magnetic field and density structure measurements using hard X-rays in a flaring coronal loop (Received July, 2008; accepted ***, 2008) |
Abstract Aims. A novel method of using hard X-rays as a diagnostic for chromospheric density and magnetic structures is developed to infer sub-arcsecond vertical variation of magnetic flux tube size and neutral gas density. |
Methods. Using Reuven Ramaty High Energy Solar Spectroscopic Imager (RHESSI) X-ray data and the newly developed X-ray visibilities forward fitting technique we find the FWHM and centroid positions of hard X-ray sources with sub-arcsecond resolution ( [MATH] ) for a solar limb flare. We show that the height variations o... |
[MATH] 150 km by mapping 18-250 keV X-ray emission of energetic electrons propagating in the loop at chromospheric heights of 400-1500 km. |
Results. Our observations suggest that the density of the neutral gas is in good agreement with hydrostatic models with a scale height of around |
[MATH] km. FWHM sizes of the X-ray sources decrease with energy suggesting the expansion (fanning out) of magnetic flux tube in the chromosphere with height. The magnetic scale height |
[MATH] is found to be of the order of 300 km and strong horizontal magnetic field is associated with noticeable flux tube expansion at a height of [MATH] 900 km. |
Key Words.: Sun:chromosphere - Sun:flares - Sun: X-rays, gamma rays - Sun:magnetic fields - Sun:activity Introduction Chromospheric magnetic fields are notoriously difficult to measure and their detailed structure is effectively inaccessible with modern observations. The difficulties of various line spectroscopic techn... |
(e.g. Solanki et al. 2006 occur because the magnetic field is relatively weak so the observed spectral lines are consequently broad and insensitive to the field. The computation of chromospheric vector magnetic field from spectral lines is also an ill-conditioned inverse problem (e.g. Metcalf et al. 1995 . In addition,... |
(Lagg 2007 . Therefore, various indirect techniques are often employed to determine the magnetic field in the chromosphere: optical observations of photospheric magnetic fields combined with extrapolation into the chromosphere |
(e.g. McClymont et al. 1997 ; radio observations of gyroresonance emission (Lang et al. 1993 ; Aschwanden et al. 1995 ; Vourlidas et al. 1997 ; White & Kundu 1997 |
The solar chromosphere being only about 2000 km thick ( [MATH] strategically covers the layer where the solar atmosphere turns from the gas-dominated lower chromosphere/photosphere into the magnetic field dominated upper chromosphere/corona. Gabriel ( 1976 has proposed that magnetic field in the chromosphere fans out (... |
Giovanelli & Jones ( 1982 have found that the canopy height should be typically 300–400 km. However, polarisation measurements by |
Landi degl’Innocenti ( 1998 suggest a very small horizontal component of the magnetic field and Schrijver & Title ( 2003 argue that the “wine-glass” shaped magnetic field should return to the photosphere near their parent flux tube. In addition, different magnetic field models predict different canopy heights |
(Solanki et al. 1999 The transport of both thermal and energetic charged particles in the solar atmosphere is governed by individual magnetic flux tubes. Therefore flare accelerated electrons one-dimensionally propagating along magnetic field lines can trace individual flux tubes from the electron acceleration site in ... |
Aschwanden et al. ( 2002 to infer the chromospheric density structure from Reuven Ramaty High Energy Solar Spectroscopic Imager RHESSI |
(Lin et al. 2002 X-ray observations. In this letter, we show that X-rays can be a diagnostic tool for the analysis of not only energetic electrons in solar flares but of the magnetic flux tubes and density structure in the chromosphere. We analyse the spatial and energy distribution of hard X-ray sources using RHESSI d... |
Electron precipitation in magnetic loops Accelerated electrons follow magnetic field lines towards the denser layers of the atmosphere, losing energy to binary collisions and scattering in angle en route. In the first instance we follow Brown et al. ( 2002 and neglect scattering. Then the flux [MATH] [electrons s -1 cm... |
at depth [MATH] from the acceleration point [MATH] , is [EQUATION] where [MATH] is the column depth [MATH] [MATH] [MATH] is the electron charge, [MATH] is the Coulomb logarithm (Brown 1973 ; Kontar et al. 2002 and we have assumed an injected flux of accelerated electrons [MATH] [electrons keV -1 cm -2 -1 ] at [MATH] . ... |
is [EQUATION] where [MATH] is the photon energy and [MATH] is isotropic bremsstrahlung cross-section (Haug 1997 Since the density [MATH] increases with [MATH] while the electron flux |
[MATH] decreases, [MATH] thus exhibits a maximum for a fixed [MATH] at a depth [MATH] that increases monotonically with [MATH] (Brown et al. 2002 . Lower energy electrons lose their energy faster, with higher energy electrons propagating deeper into the atmosphere and the higher energy X-ray emission should come from l... |
If the magnetic field strength [MATH] increases with depth (Figure ), the cross-sectional area [MATH] of the flux tube will obey the principle of magnetic flux conservation [MATH] |
(e.g. Melrose & White 1979 . The size of the X-ray source at energy [MATH] becomes a measure of [MATH] at [MATH] for limb events (Figure ), which are unaffected by albedo effects (Kontar et al. 2006 We assume in the first instance that the electrons move parallel to the field so they experience no mirror force although... |
(Brown 1972 ; Leach & Petrosian 1981 ; MacKinnon & Craig 1991 RHESSI X-ray observations We selected a large GOES M6 class X-ray flare that appeared at the eastern limb on January 06, 2004 with a hard X-ray peak at [MATH] 06:23UT. The flare has an extended coronal source visible in soft X-rays and two hard X-ray footpoi... |
[MATH] 120 keV. This flare seems ideal for our analysis since it is a limb event with one dominant source of hard X-ray emission seen in images up to |
[MATH] 300 keV. In addition, the spatially integrated count rate suggests emission above [MATH] keV. The spatially integrated photon spectrum for the same time interval (06:22:20-06:23:00 UT) has been fitted using isothermal plus thick-target model (Figure 3). The X-ray emission above 18 keV (Figure 3) is dominated by ... |
X-ray visibilities (2-dimensional spatial Fourier components) (see Schmahl et al. 2007 in ten different hard X-ray energy ranges from 18 to 250 keV have been forward fitted using single circular Gaussian source (Figure 3). Although the figure shows the fitted source matching the imaged footpoint location, the fit has b... |
[MATH] we can readily find the height of the X-ray source [MATH] , where [MATH] is radial distance of the bottom of the loop. The typical uncertainties on radial distance measurement [MATH] are around [MATH] |
[MATH] km. [MATH] is poorly known (Aschwanden et al. 2002 but crucial for the detailed analysis, therefore we incorporate it as a fit parameter. Assuming the neutral hydrogen density profile to be |
[EQUATION] where [MATH] is the density scale height, and [MATH] is the number density at height [MATH] . Various chromospheric models A-F from |
Vernazza et al. ( 1981 provide slightly different values of [MATH] though most are very close to [MATH] cm -3 . Therefore we use this [MATH] |
as a boundary condition to find the characteristic scale-height as well as [MATH] . Using equation ( ) to fit the maximum of flux spectrum given by Equation ( ) with [MATH] |
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