adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016A&A...587A..61C (fulltext)
is: (1) ID number used for cross-reference with the web version ofthis catalogue; (2) Name of the system and optical counterpart when known; (3) Peak X-ray flux in erg s cm, standardized to the 210 keV band. To do so, we begin with the X-ray flux that has been published in the literature (or in archive). We assume a power-law spectrum with a photon index ( Belloni et al. 2011 ) and the total neutral Galactic Hydrogen column density (), published by Kalberla et al. (2005) . If there is a measured published in literature derived from direct X-ray spectral analysis, we use this instead of the radio-derived interstellar one; (45) Optical or IR magnitude in the peak of the outburst and quiescence, respectively, in the AB system. To document the original observed band, we provide the name of the band in its original system; (6) Optical Galactic extinction reported in the literature. If unknown, we list the total Galactic line-of-sight absorption given by the Schlafly Finkbeiner (2011) dust maps; (7) Reported or estimated orbital period of the binary in hours; (8) References for all the parameters above. Table A.3 lists the optical/NIR photometry of the dynamical BHTs in quiescence: (1) Preferred name of the dynamically confirmed BHT; (29) Quiescent magnitudes of the dynamically confirmed BHTs. All the magnitudes were transformed to the AB system using the transformation coefficients taken from Table 2 in Frei Gunn (1994) and Eqs. (5) in Blanton et al. (2005) . To document the original observed band, we provide the name in its original system; (10) References of the magnitudes. Table A.4 provides the dynamical parameters of the BHTs: (1) Preferred name of the dynamically confirmed BHT; (2) Spectral type of the companion star; (3) The orbital period of the binary in hours; (4) The radial velocity of the companion star () in ; (5) The mass function of the BH in ; (6) The mass of the BH () in . If there is an uncertain value, we prefer to show a range of masses; (7) The binary mass ratio ; (8) The inclination of the system () in degrees; (9) The rotational broadening () in ; (10) References of all the parameters. At the bottom of each table, we present a detailed list with the particularities marked in Table A.1 and Table A.2 . This census is the most updated available and, compared to previous catalogues (e.g. Gottwald et al. 1991 ; Chen et al. 1997 ; Remillard McClintock 2006 ; Ritter Kolb 2003 ; Liu et al. 2006 , 2007 ), represents a substantial improvement in both statistics and observational parameters. While it only focuses on the population of BHs, it provides a thorough and more complete coverage of optical and NIR data and dynamical parameters 2 . 3. Analysis of the spatial distribution of BHTs Table A.1 allows us to perform a statistical study of the distribution of BHTs in the Galaxy. In Fig. 2 , we plot the 35 objects with estimated distances as viewed from the pole of the Milky Way. Filled orange circles represent the
Paper: 2020MNRAS.492.5865C (acknowledgments)
ACKNOWLEDGEMENTS We are grateful to the referee for helpful comments on the manuscript. This work was supported by a grant from the Research Grants Council of the Hong Kong Special Administrative Region, China (Project No. EdUHK 28300518) and the grants from the Education University of Hong Kong (RG2/2019-2020R and RG7/2019-2020R).
Paper: 2018A&A...614A..63S (acknowledgments)
Acknowledgements This work was supported by Sonderforschungsbereich SFB 881 “The Milky Way System” (subprojects A6 and A5) of the German Research Foundation (DFG). Funding for RAVE has been provided by: the Australian Astronomical Observatory; the Leibniz-Institut fuer Astrophysik Potsdam (AIP); the Australian National University; the Australian Research Council; the French National Research Agency; the German Research Foundation (SPP 1177 and SFB 881); the European Research Council (ERC-StG 240271 Galactica); the Istituto Nazionale di Astrofisica at Padova; The Johns Hopkins University; the National Science Foundation of the USA (AST-0908326); the W. M. Keck foundation; the Macquarie University; the Netherlands Research School for Astronomy; the Natural Sciences and Engineering Research Council of Canada; the Slovenian Research Agency (P1-0188); the Swiss National Science Foundation; the Science Technology Facilities Council of the UK; Opticon; Strasbourg Observatory; and the Universities of Groningen, Heidelberg and Sydney. The RAVE web site is at https://www.rave-survey.org . Funding for SDSS-I and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org . UCAC5 proper motions and the distances from McMillan et al. (2018) were obtained with the use of the European Space Agency (ESA) mission Gaia ( https://www.cosmos.esa.int/gaia ), processed by the Gaia Data Processing and Analysis Consortium (DPAC, https://www.cosmos.esa.int/web/gaia/dpac/consortium ). Funding for the DPAC has been provided by national institutions, in particular the institutions participating in the Gaia Multilateral Agreement. The authors are very grateful to Young-Sun Lee and Timothy C. Beers for providing their SEGUE data sample for our analysis and for fruitful discussions. We also thank the anonymous referee for the detailed suggestions, which improved the paper significantly.
Paper: 2021MNRAS.502.5668V (fulltext)
short ones (lasting a few days) and superoutbursts, which can last two weeks or longer in their rather bright ‘plateau’ phase. Typical superoutburst cycle lengths of these ‘ordinary’ SU UMa stars range from 100–500 d. Two additional subgroups can be distinguished: WZ Sge stars have longer superoutburst cycle lengths (several thousand days) and lack the normal short outbursts, while ER UMa-type dwarf novae are characterized by very short superoutburst cycles (20–40 d), and even more frequent normal outbursts in between. More than four decades ago, periodic ‘superhumps’ were detected in the SU UMa star VW Hyi (Vogt 1974 ; Warner 1975 ) during a plateau phase; later they turned out to be a characteristic feature of all SU UMa stars, but they also appear occasionally in other CV types, as classical nova remnants and nova-like stars (Patterson 1999 ). Since then, and till now, most observational efforts on SU UMa stars have been concentrated on the detection of superhumps and the determination of their properties (Kato et al. 2017 and the previous papers of that series), widely neglecting the systematic coverage of the general outburst behaviour of these and other dwarf novae. The semiregular outburst behaviour of DNe has attracted the interest of amateur astronomer societies for a long time, such as the American Association of Variable Star Observers (AAVSO), the British Astronomical Association (BAA), and the Royal Astronomical Society of New Zealand, Variable Star Section (RASNZ, VSS), and very valuable long-term light-curve studies have already been published decades ago, containing catalogues of individual outbursts with their main properties as maximum brightness, amplitude, width, and classification according to light-curve shape, as well as some statistical analyses of these properties. As typical examples, we can mention here extensive studies on SS Cyg (Mattei, Saladyga Waagen 1985 ), U Gem (Mayall 1957 ; Mattei et al. 1987 ), and VW Hyi (Bateson 1977 ). To our knowledge, no long-term studies like those of the 1960s and 1970s have been published more recently. Simultaneously, there was and still is a boom in the number of detections of new CVs, especially dwarf novae. The VSX AAVSO catalogue (Watson 2006 ) lists more than 9000 DNe, but does not give any cycle length information. The latest online version of the CV catalogue of Ritter Kolb ( 2003 , version 7.24, 2017 February) list a total of 852 DNe with known orbital period, 631 of them (74 ) classified as SU UMa, ER UMa, or WZ Sge stars. Only for 121 (19 ) are some cycle length data given. Discounting the doubtful cases, Ritter Kolb ( 2003 ) give, for a total of only 73 cases (12 ), reliable information on the lengths of the superoutburst cycle (hereafter ‘supercycle length’). However, no references are given for the cycle lengths, and it is often very difficult and time-consuming to check these data. Therefore, we here determine supercycle length values of other SU UMa stars, for which
Paper: 2015A&A...577A..89B (acknowledgments)
The calculations were carried out in the CFCA cluster of National Astronomical Observatory of Japan. We thank Andrea Mignone and the PLUTO team for the possibility to use the PLUTO code. We especially thank to Claudio Zanni for technical support. The visualization of the results used the VisIt package. V.B.-R. acknowledges support by the Spanish Ministerio de Economía y Competitividad (MINECO) under grants AYA2013-47447-C3-1-P. V.B.-R. also acknowledges financial support from MINECO and European Social Funds through a Ramón y Cajal fellowship. This research has been supported by the Marie Curie Career Integration Grant 321520. B.M.V. acknowledges partial support by the JSPS (Japan Society for the Promotion of Science): No. 2503786, 25610056, 26287056, 26800159. B.M.V. also acknowledges MEXT (Ministry of Education, Culture, Sports, Science and Technology): No. 26105521 and RFBR grant 12-02-01336-a. M.P. is a member of the working team of projects AYA2013-40979-P and AYA2013-48226-C3-2-P, funded by the Spanish Ministerio de Economía y Competitividad.
Paper: 2020MNRAS.495.4028C (fulltext)
envelope for the AGB case, apparently because outward energy transport by spiral shocks is impeded by a dense inner envelope in the RGB case. For the volume-integrated energy terms (normalized by the total initial energy), there was a high level of agreement between the two runs, but with a few notable differences. The potential energy term involving the primary core and envelope is relatively more important in the AGB case (higher particle mass, more centrally condensed envelope), while the potential energy terms involving the secondary and envelope and self-gravity of the envelope are relatively more important in the RGB case (larger envelope mass). Our setup can eventually be improved by using a synchronously rotating primary initialized at the Roche limit separation (MacLeod, Ostriker Stone 2018 ; Reichardt et al. 2019 ), although this comes at a much higher computational cost. Radiative transfer and a more realistic non-ideal equation of state are needed. Including these ingredients will allow us to model ionization and recombination, radiative cooling, and convection, all of which are likely important for the budget and redistribution of energy in the envelope, and more accurate modelling of envelope unbinding. Despite these caveats, the projection of to unbind the envelope suggests that the universal failure in previous simulations to eject it without recombination is unlikely a consequence of insufficient physics, but the result of a combination of numerical limitations and insufficient run time. Nevertheless, the need to improve both the physics and the numerical capabilities of simulations remains, while also expanding into new regions of phenomenologically relevant parameter space. APPENDIX A: INITIAL CONDITIONS Initial profiles for the AGB primary star used are presented in Fig. A1 . We refer the reader to Chamandy et al. ( 2019a ) for the same profiles for the RGB star used. Figure A1. Top panel: radial profile of gas density for the AGB star modelled using mesa (thick orange curve) and the envelope of our 3D AGB star in the simulation at t = 0 (thin black curve). Middle panel: comparison of the internal energy density profiles in the mesa model, the simulation initial envelope, and the mesa model with the equation of state replaced with an ideal gas equation of state, as in the simulation (dashed red curve). Bottom panel: comparison of (negative of) potential energy density profiles in the mesa model and the simulation initial envelope. APPENDIX B: ENERGY NON-CONSERVATION The degree to which energy conservation is satisfied for our AGB run (denoted Model A) and two lower resolution AGB runs (Models B and C) is shown in Fig. B1 , where the fractional change in the total energy (accounting for flux through the domain boundaries) is plotted. The level of adherence to energy conservation is sensitive to the resolution around the primary core particle, and, after the first periastron passage, also to the resolution around the
Paper: 2017ApJ...844..157L (fulltext)
the UDGs: to −15.5 mag, and M * = 5 × 10 7 to 5 × 10 8 M ⊙ . The slanted solid lines represent the varying surface brightness, = 23, 24, 25, 26, and 27 mag arcsec −2 (from left to right). For given surface brightness, the larger the UDGs are, the brighter (more massive) they are. The stellar masses of the UDGs are much lower than those of the bright red sequence galaxies. Figure 6. Left and middle panels: circularized effective radii vs. r ′-band absolute magnitudes (lower x -axes) and stellar masses (upper x -axes) of the UDGs and other galaxies in Abell S1063, Abell 2744, and the parallel fields. Symbols are the same as in Figure 5 . Solid lines represent isosurface brightness magnitudes of = 23, 24, 25, 26, and 27 mag arcsec −2 from left to right. Right panels: red line histograms are for the UDGs and LSB dwarfs in Abell S1063 and Abell 2744, and gray line histograms are for the UDGs in Coma. The numbers and dotted lines represent mean values of Sérsic index n for UDGs. In the right panels of Figure 6 we plotted the histograms of effective radii for Abell S1063 and Abell 2744 in comparison with that of the Coma UDGs (Yagi et al. 2016 ). The mean values, kpc for Abell S1063 and kpc for Abell 2744, are similar to that of the MW-sized UDGs in Coma, kpc. The histograms of the combined sample of UDGs and LSB dwarfs are similar to those of the large UDGs in Coma. 3.5. Spatial Distributions of the UDGs Figure 7 (left panels) displays the spatial distribution of the UDGs, LSB dwarf galaxies, and bright red sequence galaxies in the Abell S1063 and Abell 2744 fields. The circles in the left panels represent the boundaries of the radial bins used for deriving radial number density profiles. We excluded the central circular region in the UDG survey because of high surface brightness of the central galaxy. Figure 7. Spatial distributions (left panels) and radial number density profiles corrected for detection completeness (right panels) of the UDGs (large red starlets), LSB dwarfs (small red starlets), and bright red sequence galaxies (blue circles) in Abell S1063 and Abell 2744. Black dots denote the compact galaxies with . The circles in the left panels represent the boundaries of the radial bins used for deriving radial number density profiles. The central gray circle denotes the region not used for the UDG survey. Yellow starlets in the right panels are for the sum of the UDGs and LSB dwarfs. Yellow and cyan regions represent the cluster and parallel fields, respectively. We derived the radial number density profiles of the galaxies, subtracting the background contribution using the data for the HXDF. As the center for the radial profiles we adopted the center of the cD galaxy in Abell S1063 and the center of CN-1 in Abell 2744. Note that Lee Jang ( 2016 ) adopted the center of CN-2 at the southeast of CN-1, which is the brightest galaxy in the HST field of Abell 2744, for deriving radial number density profiles of the point sources (mostly
Paper: 2019MNRAS.482.4726N (acknowledgments)
ACKNOWLEDGEMENTS I acknowledge the support from the European Research Council (ERC) Consolidator Grant funding scheme ( project STARKEY , grant agreement number 615604) and from the Centre National d’Études Spatiales (CNES). I thank all the members of the STARKEY group and Professor Alessandro Bressan for the fruitful discussions. I thank the anonymous referee for carefully reading the manuscript and for the useful suggestions that improved this work. This work has made use of data from the European Space Agency (ESA) mission Gaia ( https://www.cosmos.esa.int/gaia ), processed by the Gaia Data Processing and Analysis Consortium (DPAC, https://www.cosmos.esa.int/web/gaia/dpac/consortium ). Funding for the DPAC has been provided by national institutions, in particular the institutions participating in the Gaia Multilateral Agreement. This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation.
Paper: 2017MNRAS.470..324T (fulltext)
± 0.10, assuming no change in flux between the two observations), similar to the compact jets of other X-ray binaries, although we cannot rule out a steep spectrum within uncertainties. We find that the spectral index measurements within the 10 GHz band on the other observing epochs are typically unconstraining (with a typical error of ±0.5, due to the faintness of the radio emission and the small lever arm in frequency), so we cannot distinguish between optically thin and thick emission. The spectral index measurements for all epochs are listed in Table 3 . 4.5 The radio–X-ray correlation The radio ( L R = 4π D 2 ν S ν , where D is the distance to the source) and X-ray luminosities of our sample are displayed in Fig. 5 individually, and together in Fig. 6 , with the addition of other black hole and neutron star systems from the literature. Figure 5. The radio and X-ray luminosities of Cen X−4, SAX J1808.4−3658 and IGR J00291+5934. Hollow symbols mark upper limits. The labelled data points of SAX J1808.4−3658 refer to the same observations in Fig. 4 . IGR J00291+5934 (and possibly SAX J1808.4−3658, but with larger scatter) seems to follow the familiar correlation seen in black holes and some atoll sources, with β = 0.77 ± 0.18 (90 per cent confidence interval, shaded area). Positions of X-ray binaries on the radio–X-ray plane are best measured from strictly simultaneous observations. Among our multiwavelength campaign, Cen X−4 has strictly simultaneous data. For the other sources, we interpolate the X-ray light curve to estimate the X-ray flux at the time of the radio observations, by fitting a piecewise linear function in the log L X /time space. The 3 σ upper limit of 14 μJy on the radio emission of Cen X−4 in quiescence only constrains its radio luminosity to being less than that of the radio-fainter black hole jets (e.g. XTE J1118+480; Gallo et al. 2014 ). The limit lies just above the proposed tMSP track, slightly below the quiescent black hole systems. In addition, we also place the 1979 outburst of Cen X−4 on the radio/X-ray plane for the first time. For the radio luminosity, we use the data reported by Hjellming et al. ( 1988 ) in two bands (at 1.49 and 4.9 GHz). For days when observations were carried out in both radio bands, we use the measurement with smaller uncertainties, with the assumption of a flat spectrum. We find that during the 1979 outburst, Cen X−4 reached radio and X-ray luminosities similar to the atoll sources. Assuming it follows a single, continuous radio–X-ray correlation from the peak of the outburst down to quiescence (like the black hole systems V404 Cyg, Corbel et al. 2008 ; and GX 339–4, Corbel et al. 2013 ), we obtain a lower limit for its correlation slope β ≳ 0.5. The upper limits for IGR J17511–3057 make it less luminous in the radio band than typical atoll sources (Fig. 6 ), but we cannot rule out radio emission at a level similar to that of EXO 1745−248, the radio-detected neutron star with the lowest
Paper: 2019MNRAS.484.3552B (acknowledgments)
ACKNOWLEDGEMENTS This publication makes use of data from the GLQ database in the GLENDAMA archive available at http://grupos.unican.es/glendama/database/ . The research activities of the Observational Astronomy Board of the Federal University of Rio Grande do Norte (UFRN) are supported by continuous grants from CNPq, CAPES, and FAPERN Brazilian agencies. We also acknowledge financial support from INCT INEspaço/CNPq/MCT. AB acknowledges a CAPES PhD fellowship. ICL acknowledges a CNPq/PDE fellowship. We warmly thank the anonymous reviewer for the fruitful comments and suggestions that greatly improved this work.
Paper: 2018MNRAS.480.4069M (acknowledgments)
ACKNOWLEDGEMENTS We would like to thank the referee for a constructive and useful comments on the manuscript. AK is very grateful to Rick Kessler and Dan Scolnic for many enlightening discussions on the cosmological analyses using supernovae type Ia samples and effects of survey selection function. AK was supported by a NASA ATP grant NNH12ZDA001N, NSF grant AST-1412107, and by the Kavli Institute for Cosmological Physics at the University of Chicago through grant PHY-1125897 and an endowment from the Kavli Foundation and its founder Fred Kavli. EJ is supported by the Argonne Leadership Computing Facility, a U.S. Department of Energy, Office of Science User Facility operated under Contract DE-AC02-06CH11357. CM and JW were supported by the U.S. Department of Energy, Office of Science, Office of Advanced Scientific Computing Research (ASCR) under Contract DE-AC02-06CH11347. JW was also supported by ASCR through award DE-SC0014205. Numerical experiments presented in this paper were carried out on the midway computing cluster maintained by the University of Chicago Research Computing Center. We acknowledge the Center and its staff for support.
Paper: 2019ApJ...887...14A (acknowledgments)
We thank Paule Sonnentrucker for useful comments and Robert Snuggs for technical assistance.