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2.2 The oscillation calculations and iron accumulation We adapted the Eggleton evolution code so that the output is suitable for pulsation calculations. In practice, this implied calculating some additional physical quantities during the evolution, and modifying the mesh to have sufficient meshpoints in the stellar env... |
Dupret ( 2001 . We determined the theoretical frequency spectrum up to [MATH] since it is expected that higher order modes are geometrically cancelled. |
Charpinet et al. ( 1996 has established that the excitation of sdB oscillatons is related to a local enrichment of iron in the stellar envelope caused by diffusion. Radiative levitation is expected to set up significant chemical gradients within a diffusion timescale of [MATH] [MATH] yr, and consequently iron accumulat... |
(Michaud et al. 1989 ; Chayer et al. 1995 . Time-dependent diffusion calculations (Fontaine et al. 2006 show that, after [MATH] [MATH] yr, many pulsation modes are excited. Since element diffusion is not treated in the evolution code, we used an approximation for the iron accumulation, assuming that the iron only affec... |
[EQUATION] with the initial condition [MATH] . The width [MATH] and accumulation timescale [MATH] yr are chosen such that iron is only increased in the region [MATH] , and [MATH] , which is loosely based on the time-dependent diffusion calculations of |
Fontaine et al. ( 2006 and the equilibrium profiles of Charpinet et al. ( 1997 . Our parametric approximation is rather ad hoc, but since we are interested in the relative differences between two different scenarios, the exact shape of the iron profile is not crucial here. We will discuss the influence of the iron prof... |
Results 3.1 Effects of the iron accumulation In Fig. , we show [MATH] throughout the star for different ages of an sdB star. Note that the temperature range [MATH] corresponds to a very narrow mass shell of [MATH] [MATH] . In Fig. c, we included the backreaction of convective mixing on our parametric iron profile durin... |
[MATH] [MATH] ) and [MATH] [MATH] ). This is caused by two narrow convective layers due to iron and helium ionization, respectively. Interestingly, we found that the convective region around [MATH] would not be present without iron accumulation. We determined that the slightly perturbed iron profile has a negligible ef... |
In Fig. , we show the effect of different iron abundance profiles on the sdB evolution in the [MATH] diagram. We compared our parametric approach (Eq. ) with the case of no iron enhancement, and a uniform enhancement of [MATH] in the whole envelope, as used in studies of mode excitation (Jeffery & Saio 2006 . It is evi... |
We also compared the effect on the excitation, and find that our approximation of the iron profile can excite almost as many modes as the [MATH] enhancement, see Fig. a-c. This is understood in terms of the driving mechanism in sdB stars, which is associated with the iron opacity bump such that accumulating iron in thi... |
3.2 The stellar models Following the procedure as described in § , we constructed a grid of canonical, i.e. post-He-flash, sdB models (hereafter called grid [MATH] ) with masses in the range [MATH] in steps of 0.01 M . The maximum mass we obtained for the degenerate core of an RGB star is [MATH] , thus we did not consi... |
[MATH] , considered are 0.0001, 0.0003, and 0.0006 M , where [MATH] is defined as the total mass of the hydrogen content directly after the removal of the envelope. Thus, we have 18 sdB evolution tracks, which we followed until the end of He-core burning. After each [MATH] yr of sdB evolution, the seismic properties we... |
The grid of non-canonical sdB stars (hereafter called grid [MATH] ) consists of 5 tracks: ( [MATH] (M ), [MATH] (M )) = (0.44, 0.005), (0.45, 0.005), (0.46, 0.005 ), (0.47, 0.0075 ), and (0.47, 0.005). Along these tracks we have in total 98 seismic models, again taken after each [MATH] |
yr of sdB evolution. In Table , more details about the models are given. The sdB evolution tracks and the seismic models can be seen in Fig. . The tracks start directly after the removal of the envelope. For the post-flash models, this corresponds to the zero-age EHB. The post-non-degenerate models have hydrogen extend... |
yr) of H-shell burning, before starting He-shell burning. 3.3 Comparing two representative models We examined the physical differences in the interior structure of a post-flash [MATH] ) and a post-non-degenerate ( [MATH] ) sdB with same [MATH] and [MATH] . We chose as representative the models circled in Fig. b at [MAT... |
Two important quantities in stellar pulsation theory are the the Brunt-Väisäla frequency [MATH] and the Lamb frequencies [MATH] [EQUATION] |
[EQUATION] [EQUATION] where [MATH] is the molecular weight and [MATH] is the adiabatic sound speed. When [MATH] , the Ledoux criterion for dynamical stability is violated (Ledoux 1947 . Thus, in Fig. b&c, the convective regions can be clearly identified. The innermost one is related to the convective core, and the oute... |
[MATH] , and the [MATH] -modes, [MATH] , where [MATH] is the angular pulsation frequency. It is apparent that [MATH] -modes are deep interior modes, while [MATH] -modes probe the superficial outer layers as pointed out by Charpinet et al. ( 2000 |
Clearly, models [MATH] and [MATH] have very different physical characteristics. To establish how this affects their seismic properties, we compare their frequencies in Fig. . Since the large frequency separation [MATH] is mainly dependent on the dynamical timescale, we see that [MATH] at high frequencies is more or les... |
[MATH] modes are not the same for these two models; for model [MATH] this range is [3.4 mHz, 10.6 mHz], and for model [MATH] , it is [5.7 mHz,14.2 mHz]. The excited modes thus have lower frequencies in model [MATH] than in model [MATH] . To understand this, we compare in Fig. h the work integral of these two models for... |
A first possible origin of the differences could come from the opacity, since the driving is a [MATH] -mechanism operating in the iron opacity bump. Fig. e shows that the opacity is slightly larger for model |
[MATH] . The driving is thus a little more efficient in model [MATH] . But this is not the main source of differences. Since the envelope H-fraction is much smaller in model [MATH] [MATH] ) than in model [MATH] [MATH] ), the molecular weight is larger and, at given temperature, the density is significantly higher ( [MA... |
[MATH] because of the higher density, as can be seen in Fig. f&g. Hence, the difference between the temperature at the last node and at the surface, [MATH] , is greater in model [MATH] than in model |
[MATH] . This is exactly what we find in Fig. i, where the eigenfunction [MATH] is given. In terms of the temperature, the last node is closer to the surface in model [MATH] than it is in model [MATH] . To get the same driving as in model [MATH] , the last node of model [MATH] would have to be deeper in the star, which... |
3.4 Comparing two grids of models We investigated if it is possible to distinguish between a post-He-flash ( [MATH] and a post-non-degenerate sdB stellar model ( [MATH] ) from observed oscillation modes. Imagine we observed the frequencies of [MATH] ; is it then possible to find an acceptable seismic match in our grid ... |
[MATH] )? We took as ‘observed’ frequencies those of unstable modes up to [MATH] . We did this for each model [MATH] in grid [MATH] , thus finding the best seismic match within grids [MATH] and [MATH] |
Since frequency separations follow from asymptotic relations for [MATH] -modes, the frequency is a natural quantity for model comparison. Despite this, periods have been used more often in the literature so far, when comparing observed modes of sdB stars with those predicted by models. We also considered period matchin... |
To quantify ‘acceptable’, we used the merit function [EQUATION] where [MATH] is one of the [MATH] excited frequencies of star [MATH] , and [MATH] is the correspondingly matched frequency of star [MATH] , expressed in mHz. The frequency matching is done such that [MATH] is minimized by brute-force fitting. It is clear t... |
(i) We are not able to identify the modes, [MATH] and [MATH] of the ‘observed’ star [MATH] are unknown, and the ‘observed’ frequencies are allowed to be matched with both stable and unstable frequencies of the ‘theoretical’ model [MATH] |
(ii) Same as (i), except the modes are identified, thus the [MATH] -value must be matched. (iii) Same as (i), except [MATH] and [MATH] are known within errors of [MATH] and [MATH] |
(iv) Same as (i), except the ‘observed’ frequencies are only matched with unstable ‘theoretical’ frequencies, i.e. assuming that the theory correctly predicts which frequencies are exited and which are not. |
In Fig. , we show [MATH] for each gridpoint in grids [MATH] and [MATH] for the scenarios (i)-(iv). The matches with low [MATH] are visible as dark diagonal regions, This is a result of the change in frequencies during the sdB evolution. From Fig. (i)-(iii), it is clear that the distinction between models [MATH] and [MA... |
The matches with lowest [MATH] are circled in Fig. and details of these models are shown in Fig. and Table . For all scenarios the same model [MATH] gives the best match, namely the last model of evolutionary track [MATH] . We understand that in terms of only the higher frequencies with radial order [MATH] being excite... |
As a comparison to a real case, we considered the optimal model for PG 0014$+$067 , for which Brassard et al. ( 2001 found [MATH] , where [MATH] is a merit function based on mode period comparison. Translated to our frequency merit function, this is equivalent to [MATH] . Although we find, in principle, seismic matches... |
Discussion & conclusions We studied the so far neglected, post-non-degenerate sdB stars and compared their physical and seismic characteristics with those of canonical post-flash sdB stars, both formed in the CEE channel. The results presented here are a first step in distinguishing these two kinds of sdB stars on the ... |
which we started in Hu et al. ( 2007 and Vučković et al. ( 2007 We find that, in principle, a post-non-degenerate sdB star may appear as an EC 14026 star with similar pulsation frequencies as the canonical post-He-flash sdB star, although it is not likely. Additional observables, such as spectroscopic [MATH] and |
[MATH] determinations and/or empirical mode identification from observables enable us to distinguish the two types of sdB stars more decisively. The frequency range of the unstable modes is also an important discriminator between the two formation channels. In general, for the same [MATH] and [MATH] values, the excited... |
Up to now, there have not been any evolutionary models of sdB stars available that include the coupling between diffusion and evolution consistently. This is a deficiency, since iron accumulation due to radiative levitation is responsible for the pulsational instability in these stars (Charpinet et al. 1996 . Also, it ... |
Fontaine et al. ( 2006 that the iron accumulation changes the frequencies significantly. In our study, we have parametrized the iron accumulation, so that we can, at least in an approximative manner, simultaneously take into account the effects of iron enhancement and evolution on the pulsation modes. |
Here we have not considered the influence of the other diffusive processes, i.e. diffusion due to gradients of pressure, temperature, and concentration. To a certain extent this can affect our results, because one of the main differences between the two types of sdB stars is the chemical composition of the stellar enve... |
(Richard et al. 2002 ; Michaud et al. 2007 . Since the envelopes of the post-non-degenerate sdB stars extend to [MATH] (i.e. [MATH] K), we do not expect diffusion to wash away all the qualitative differences in the chemical profiles, although the differences may be less pronounced. Diffusion, however, will significantl... |
We have made a modest grid of models that is sufficient for our comparative study. Detailed seismic modelling of an observed star, however, will require a finer grid. For now, we have chosen not to make sdB models above 0.47 , since this is the maximum mass the degenerate He-core of a red giant with [MATH] can have bef... |
[MATH] allows the He-core to grow up to 0.48 M on the RGB. However, we find that, in order to excite modes in these low metallicity stars, an iron enhancement greater than a factor 10 is required. This was to be expected, since Charpinet et al. ( 1996 found unstable pulsation modes for models with uniform [MATH] in the... |
In this paper, we have focused on the short-period [MATH] -mode sdB pulsators. The case of the long-period [MATH] -mode sdB pulsators is, although challenging from an observational point of view, an additional very interesting theoretical case study. The [MATH] -modes only probe the outermost layers, and hence are less... |
Acknowledgements. We are grateful to P. P. Eggleton for the use of his evolution code, and to E. Glebbeek and S. de Mink for their help with this code. We thank W. van Ham for his help with the frequency-matching algorithm. We would also like to thank M. Vučković, R. Østensen, and M. D. Reed for stimulating discussions... |
# Source: arxiv 0808.3085 # Title: Monitoring Supergiant Fast X-ray Transients with Swift. III. Outbursts of the prototypical SFXTs IGR J17544-2619 and XTE J1739-302 # Sections: all # Downloaded: 2026-03-02T07:58:28.520024+00:00 |
Monitoring Supergiant Fast X-ray Transients with Swift. III. Outbursts of the prototypical SFXTs IGR J17544 [MATH] 2619 and XTE J1739 [MATH] 302. |
Abstract IGR J17544 [MATH] 2619 and XTE J1739 [MATH] 302 are considered the prototypical sources of the new class of High Mass X–ray Binaries, the Supergiant Fast X–ray Transients (SFXTs). These sources were observed during bright outbursts on 2008 March 31 and 2008 April 8, respectively, thanks to an on-going monitori... |
Subject headings: X-rays: individual: (IGR J17544 [MATH] 2619, XTE J1739 [MATH] 302) 1. Introduction IGR J17544 [MATH] 2619 and XTE J1739 [MATH] 302 are confirmed members of the new sub-class of High Mass X–ray Binaries, the Supergiant Fast X–ray Transients (SFXTs), whose members have been maily discovered with the INT... |
(Sidoli et al., 2008b The spectral properties are reminiscent of those of accreting pulsars, thus it is likely that several members of the class are actually hosting neutron stars, although the spin period has been measured only in two SFXTs (AX J1841.0–0536, Bamba et al. 2001 IGR J11215–5952, Swank et al. 2007 ). |
IGR J17544 [MATH] 2619 was discovered (Sunyaev et al., 2003 with IBIS/ISGRI on-board INTEGRAL on 2003 September 17 during a 2 hour flare reaching 160 mCrab (18–25 keV). During a [MATH] observation, both the quiescence level and the onset of an outburst was caught (in’t Zand, 2005 , observing a dynamic range as large as... |
located at 3.6 kpc (Rahoui et al., 2008 Several other bright flares have been observed with INTEGRAL in 2003, 2004 and 2005 ( Grebenev et al. 2003 Grebenev et al. 2004 |
Sguera et al. 2006 Walter & Zurita Heras 2007 Kuulkers et al. 2007b ), with flare durations ranging from 0.5 to about 10 hours, reaching peak fluxes of 400 mCrab (20–40 keV). More recently, two new outbursts were detected with the Swift satellite, on 2007 November 8 (Krimm et al., 2007 |
and on 2008 March 31 (Sidoli et al., 2008a , 144 days apart. The flux at peak observed with Swift /BAT was 165 mCrab (20–40 keV). The source was also active on 2007 September 21, with a fainter flaring emission up to 30–40 mCrab (20–60 keV), as observed with IBIS/ISGRI on-board INTEGRAL |
(Kuulkers et al., 2007a XTE J1739 [MATH] 302 was discovered with [MATH] after a short outburst in August 1997 (Smith et al., 1998 , and displayed a spectrum well fitted with a bremsstrahlung model with a temperature of [MATH] 22 keV, reaching a peak flux of 3.6 [MATH] 10 -9 erg cm -2 -1 (2–25 keV). Later, several other... |
located at 2.7 kpc (Rahoui et al., 2008 Upper limits to the quiescent emission were placed with ASCA observations (Sakano et al., 2002 at a level of [MATH] 1.1 [MATH] 10 -12 erg cm -2 -1 Bright outbursts (up to 300 mCrab) were detected with IBIS/ISGRI in 2003 March, and 2004 March (Sguera et al., 2006 Frequent flaring ... |
with the Swift X-ray Telescope (XRT). This outburst was also observed by the INTEGRAL/JEM-X monitor, which detected a flare starting 5 hours before the flares seen with |
Swift (Chenevez et al., 2008 Here we report on the detailed analysis of the [MATH] data of two recent outbursts from these two prototypical SFXTs: the bright flares that occurred on 2008 March 31 (Sidoli et al., 2008a from IGR J17544 [MATH] 2619 and on 2008 April 8 from XTE J1739 [MATH] 302 (Romano et al., 2008a These ... |
2. Observations and data analysis IGR J17544 [MATH] 2619 and XTE J1739 [MATH] 302 triggered the Swift /BAT on 2008 March 31 20:50:47 UT (image trigger=308224, Sidoli et al., 2008a and on 2008 April 8 21:28:15 UT (image trigger=308797, Romano et al., 2008a respectively. In both occasions, Swift slewed to the target, all... |
Swift /BAT (but not reported in literature) on 2007 September 29 and October 4 (both 8- [MATH] , 100 mCrab). There were smaller flares in the BAT both before and after the 2008 April 8 trigger from XTE J1739 [MATH] 302, with some spikes reaching [MATH] mCrab starting at 2008 April 8 16:52 UT and continuing up to 2008 A... |
The BAT data were analysed using the standard BAT software within FTOOLS ( Heasoft , v.6.4). Mask-tagged BAT light curves were created in the standard 4 energy bands, 15–25, 25–50, 50–100, 100–150 keV, and rebinned to achieve a signal-to-noise ratio (S/N) of 5. BAT mask-weighted spectra were extracted over the time int... |
using the latest spectral redistribution matrices. For our spectral fitting (XSPEC v11.3.2) we applied an energy-dependent systematic error vector |
The XRT data were processed with standard procedures ( xrtpipeline v0.11.6), filtering, and screening criteria by using FTOOLS. We considered both WT and PC data, and selected event grades 0–2 and 0–12, respectively. When appropriate, we corrected for pile-up. To account for the background, we also extracted events wit... |
Throughout this paper the uncertainties are given at 90% confidence level for one interesting parameter unless otherwise stated. When fitting the broad band spectra during the two bright flares, we included factors in the spectral fitting to allow for normalization uncertainties between the two instruments. The constan... |
3. Results 3.1. Light curves Figure shows the Swift /XRT 0.2–10 keV light curve of IGR J17544 [MATH] 2619 and XTE J1739 [MATH] 302 throughout our 2008 monitoring program, background-subtracted and corrected for pile-up, PSF losses, and vignetting. All data in one segment were generally grouped in one point (with the ex... |
Figure and show the detailed light curves in several energy bands during the brightest part of the two outbursts, together with the 4–10/0.3–4 keV, 25–50/15–25 keV hardness ratios. Fitting the IGR J17544 [MATH] 2619 4–10/0.3–4 keV hardness ratio as a function of time to a constant model yields a value of [MATH] |
and [MATH] for 30 degrees of freedom (d.o.f.). For XTE J1739 [MATH] 302 we obtain a value of [MATH] and [MATH] for 21 degrees of freedom (d.o.f.). |
3.2. Spectroscopy of IGR J17544 [MATH] 2619 The XRT/WT spectrum, extracted during the peak of the outburst (observation 00308224000, see Table ), results in a quite hard X–ray emission. Adopting an absorbed power law, we obtain a photon index of |
[MATH] , and a high column density, [MATH] cm -2 [MATH] for 143 d.o.f.). The unabsorbed flux in the 2–10 keV band is [MATH] erg cm -2 -1 A 3- [MATH] upper limit to the equivalent width of an iron line at 6.7 keV can be placed at 62 eV. A contour plot is shown in Fig. for the single power-law model fit to the WT spectru... |
[MATH] cm -2 (C-stat [MATH] for 63.24% of 10 Monte Carlo realizations with statistics [MATH] C-stat). The unabsorbed flux in the 2–10 keV band is [MATH] erg cm -2 -1 A summary of the model parameters can be found in Table This table also lists, for comparison, the spectral parameters obtained from other XRT observation... |
We fit the simultaneous BAT [MATH] XRT spectra in the time interval 168–475 s since the BAT trigger. Several models typically used to describe the X–ray emission from accreting pulsars in HMXBs were adopted (White et al., 1983 For the spectral fitting we considered BAT counts up to 50 keV (above this energy the statist... |
Masetti et al. 2006 Ferrigno et al. 2008 ). Adopting this more physical description of the spectrum, a Comptonization model ( compTT in XSPEC, |
Titarchuk 1994 ), we obtain a cold plasma (we assumed a spherical geometry for the Comptonization plasma) with a well constrained temperature of [MATH] 4 keV, and an optical depth of 19 [MATH] |
3.3. Spectroscopy of XTE J1739 [MATH] 302 The XRT/WT spectrum (observation 00308797000) extracted during the early part of the outburst was fit with an absorbed power law, obtaining a photon index of |
[MATH] , and a high column density, [MATH] =( [MATH] cm -2 [MATH] for 35 d.o.f.). The unabsorbed flux in the 2–10 keV band is [MATH] erg cm -2 -1 A contour plot is shown in Fig. for the single power-law model fit to the WT spectrum, compared with out-of-outburst emission (Sidoli et al., 2008b The PC data of the same se... |
[MATH] for 24 d.o.f.), and an unabsorbed flux in the 2–10 keV band of [MATH] erg cm -2 -1 The model parameters are summarized in Table |
Similarly to the procedure we adopted for IGR J17544 [MATH] 2619 we fit the simultaneous XRT [MATH] BAT spectra of XTE J1739 [MATH] 302 in the 0.3–10 keV and the 14–60 keV energy bands, respectively. Adopting typical models used to describe the X–ray emission from HMXBs, as in the case of IGR J17544 [MATH] 2619, we obt... |
[MATH] =( [MATH] cm -2 , although significantly better fits are obtained with a cut-off at high energies. All models (powerlaw with cutoff or Comptonizing plasma model) result in equally satisfactory deconvolutions of the 0.3–60 keV emission. In Fig. we show the result obtained adopting a power-law with a high energy c... |
4. Discussion Here we report on Swift observations of IGR J17544 [MATH] 2619 and XTE J1739 [MATH] 302 during two bright flares, observed for the first time simultaneously in a broad energy range, from 0.3 to 50–60 keV (the highest energy where the spectroscopy is meaningful with BAT in these two sources). Indeed, befor... |
Sguera et al. 2006 or at softer energy bands Smith et al. 2006 in’t Zand 2005 ). The only SFXT previously observed simultaneously in a wide X–ray band was IGR J16479–4514, during a flare caught with the Swift satellite (Romano et al., 2008b |
The X–ray spectroscopy shows that these two SFXTs, which are considered the prototypes of this new class of HMXBs, have different properties during the bright flares. IGR J17544 [MATH] 2619 is one order of magnitude less absorbed than XTE J1739 [MATH] 302, and displays a significantly flatter spectrum below 10 keV, wit... |
(in’t Zand, 2005 , where the absorbed powerlaw fit resulted in a photon index of 0.73 [MATH] , a column density of (1.36 [MATH] [MATH] 10 22 cm -2 , and a peak flux of [MATH] [MATH] 10 -9 erg cm -2 -1 |
The broad band analysis shows that IGR J17544 [MATH] 2619 displays a quite sharp cutoff at 18 [MATH] keV (when using the power law model with a high energy cut-off, highecut in Table ) or a well constrained temperature for the Comptonizing electrons (in the comptt model in XSPEC) at 4–5 keV. Instead, in XTE J1739 [MATH... |
The observations we are reporting here are part of an on-going monitoring campaign of four SFXTs with Swift (Sidoli et al., 2008b , which started on 2007 October 26. The two bright flares discussed here are the first from these two SFXTs, since the start of the campaign, which could be simultaneously covered with both ... |
Regarding the 0.3–10 keV spectra (fitted with a simple absorbed power-law), XTE J1739 [MATH] 302 appears to be much more absorbed during the flare than during the out-of-outbust emission (see Fig. ), while the photon index is similar, within the large uncertainties (Sidoli et al., 2008b Similar changes in the absorbing... |
Instead, IGR J17544 [MATH] 2619 shows a significantly harder spectrum during the flare, and a lower column density than during the out-of-outburst phase reported in Sidoli et al. ( 2008b , obtained summing together all the XRT data available from 2007 October 27 to 2008 February 28. During the out-of-outburst phase, th... |
(Sidoli et al., 2008b We also compared the bright flare spectroscopy with the spectrum extracted from one of the observations obtained a few days before the bright flare from IGR J17544 [MATH] 2619 (dashed contours in Fig. ). A hardening of the IGR J17544 [MATH] 2619 spectrum during the flaring activity is evident. A s... |
Different mechanisms have been proposed to explain the bright and short duration flaring activity in this new class of sources. Some models are related to the structure of the supergiant companion wind, involving spherically simmetric clumpy winds (see e.g. in’t Zand 2005 Negueruela et al. 2008 or anisotropic winds (Si... |
In Fig. we compare the light curves during bright flares from four SFXTs, all observed with Swift the two reported here from IGR J17544 [MATH] 2619 and XTE J1739 [MATH] 302, together with the one observed from IGR J11215–5952 (Romano et al., 2007 and from IGR J16479-4514 (Romano et al., 2008b All light curves during br... |
The wide band spectra during outbursts display high energy cut-offs (assuming the model with a power-law modified at high energy by a cutoft, highecut in XSPEC), although differently constrained in the two sources: in IGR J17544 [MATH] 2619 it is at 18 [MATH] keV, in XTE J1739 [MATH] 302 it lies below 13 keV. These cut... |
and IGR J16479–4514 (Romano et al., 2008b Facilities: Swift We thank the [MATH] team for making these observations possible, in particular Scott Barthelmy (for his invaluable help with BAT), the duty scientists, and science planners. PR thanks INAF-IASF Milano, where part of the work was carried out, for their kind hos... |
# Source: arxiv 0808.3117 # Title: The Photometric Variability of HH 30 # Sections: all # Downloaded: 2026-03-02T07:58:29.851980+00:00 |
The Photometric Variability of HH 30 Abstract HH 30 is an edge-on disk around a young stellar object. Previous imaging with the Hubble Space Telescope has show morphological variability that is possibly related to the rotation of the star or the disk. We report the results of two terrestrial observing campaigns to moni... |
HH 30 es un disco visto casi de canto alrededor de un objeto estelar joven. Imágenes previas del Hubble Space Telescope muestran una variabilidad morfológica que posiblemente esté relacionada con la rotación de la estrella o el disco. Reportamos los resultados de dos campañas observacionales realizadas con un telescopi... |
\addkeyword accretion, accretion disks \addkeyword circumstellar matter \addkeyword stars: individual (HH 30) \addkeyword stars: pre-main sequence |
0.1 Introduction High-resolution images show that HH 30 is a compact bipolar reflection nebula bisected by a dark lane (Burrows et al. 1996). Its location in the L1551 molecular cloud and similarity to the model images of Whitney & Hartman (1992) led immediately to the conclusion that HH 30 is an optically-thick circum... |
An interesting aspect of HH 30 is the prominent morphological variability (Burrows et al. 1996; Stapelfeldt et al. 1999; Cotera et al. 2001; Watson & Stapelfeldt 2007). This variability includes changes in the contrast between the brighter and fainter nebulae over a range of more than one magnitude; changes in the late... |
Two mechanisms have been suggested for the lighthouse. Wood & Whitney (1998) suggested non-axisymmetric stellar accretion hot-spots. Stapelfeldt et al. (1999) suggested voids or clumps in the inner disk. AA Tau seems to be a prototype for both mechanisms, apparently possessing both inclined hot spots and occulting inne... |
The two mechanisms are likely to be periodic, as they are expected to be tied to stellar rotation and orbital motions. Therefore, there is a hope that we might see a periodic modulation in the integrated photometry of HH 30, as one might expect the nebulae to be observed to be brighter when the lighthouse beam is point... |
0.2 Observations 0.2.1 Data Set 1 We observed HH 30 with the 84 centimeter telescope of the Observatorio Astronómico Nacional on Sierra San Pedro Mártir on 24 of the 28 nights between 1999 January 29 and 1999 February 25. We used the SITe1 [MATH] CCD binned [MATH] with the observatory’s [MATH] (2 mm KG3 and 2 mm OG570)... |
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