adsabs/WIESP2022-NER huggingface dataset:
Paper: 2017A&A...604A..48S (fulltext)
of a WD (e.g. Shen Bildsten 2007 ; Nomoto et al. 2007 ). The presence of the shell burning prior to the outburst was also suggested by Ramsay et al. (2016) . However, the high accretion rate requires a high mass transfer ratio (i.e. the mass accretion rate divided by the mass-loss rate from the donor). A very efficient wind mass transfer was suggested, for example, by Mohamed Podsiadlowski (2012) , for symbiotic stars with slow and dense winds. Skopal Cariková (2015) suggested that the high wind-mass-transfer efficiency in symbiotic stars can be caused by compression of the wind from their slowly rotating giants. Direct indication of wind focusing was found by Shagatova et al. (2016) for the symbiotic binary SY Mus. To increase the luminosity of the burning WD by a factor of 10, to values observed around the maximum of the 2015 outburst, a transient increase of the accretion rate to is needed to generate the luminosity due to hydrogen burning, where erg g is the energy production of 1 gram of hydrogen due to the nuclear fusion of four protons, and is the hydrogen mass fraction in the accreted matter. The required exceeds the stable-burning limit, which leads to blowing optically thick wind from the WD ( Hachisu et al. 1996 ). The enhanced wind is ionised by the hot WD’s pseudophotosphere, and thus converts a fraction of its stellar radiation to the nebular emission. The corresponding increase of the EM (Fig. 4 ) then causes a relevant brightening in the LC, which we indicate as the outburst (Fig. 2 ). However, the principal question remains as to the origin of this material. Probably, the red giant is responsible for a transient increase of the mass transfer ratio, which temporarily enhances . In general, Bisikalo et al. (2006) found that variations in the velocity of the wind from the giant can cause a disruption of the accretion disk, leading to the infall of a considerable amount of matter ( ) onto the WD surface, thus inducing the outburst. For Z And, Sokoloski et al. (2006) suggested that its 2000-02 outburst was triggered by the influx of hydrogen-rich material from a dwarf-nova-like disk instability at its start that consequently enhanced nuclear shell burning on the WD. 5. Summary We analysed primarily high cadence optical spectroscopy and multicolour photometry obtained prior to and during the first Z And-type outburst of the symbiotic nova AG Peg that began in June of 2015. We determined the fundamental parameters , and of the WD’s pseudophotosphere and its mass-loss rate, , in the course of the outburst (Fig. 7 ). Monitoring the evolution of the He ii 4686 Å, H and the Raman-scattered 6825 Å line profiles allowed us to suggest the ionisation structure of the active hot component in AG Peg (Fig. 12 ) and to outline the nature of its outburst. The main results of our analysis may be summarised as follows. 1. The nova-like outburst of AG Peg terminatedaround 1997 when a wave-like orbitally-related variation in theLC reached its
Paper: 2018ApJ...854...84A (acknowledgments)
We thank James Stone for helpful discussions and comments, as well as for making Athena++ publicly available. Financial support for this study has been provided by NSF via grants AST-1409524 and AST-1515763, and by NASA via grants 14-ATP14-0059 and 15-XRP15-2-0139. All simulations were carried out using facilities supported by the Princeton Institute of Computational Science and Engineering.
Paper: 2017ApJ...845..171B (acknowledgments)
This research has made use of NASA’s Astrophysical Data System. We would like to thank the Space Telescope Science Institute (STScI) and High Point University (HPU) for facilitating this research. Some/all of the data presented in this paper were obtained from the Mikulski Archive for Space Telescopes (MAST). STScI is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. Support for MAST for non- HST data is provided by the NASA Office of Space Science via grant NNX09AF08G and by other grants and contracts. Software: gPhoton (Million et al. 2016a , 2016b ), FaRVaE (Boudreaux 2017 ), SExtractor (Bertin Arnouts 1996 ), Scipy (Oliphant 2007 ; Millman Aivazis 2011 ), SAOImage ds9.
Paper: 2016A&A...588A..66T (fulltext)
overall He spread in 47Tuc is small, a small population of extreme chemistry stars may exist, which is also supported by the fact that only a handful of stars with such strong oxygen depletion has been found by e.g. Cordero et al. (2014) . The AGB polluters in Ventura et al. (2014) are in a mass range where the stars experience a mild Hot Bottom Burning, which is where the Mg-Al burning chain is activated. Their models predict an enhancement of the AGB surface abundances of Al of up to +0.5 dex, whereas the overall Mg abundances are only barely touched, with a depletion of at most 0.04 dex. Considering that the pure AGB ejecta need to be diluted with gas of pristine composition, no detectable variation in the Mg abundance is expected, consistent with our results. On the other hand, one would expect a small variation in [Al/Fe] in the polluted stars, which has indeed been reported by both Carretta et al. (2013) and Cordero et al. (2014) . That we do not see any statistically significant variation in [Al/Fe] in our sample of stars is likely a consequence of our small sample size, as discussed in Paper I. In Fig. 19 , we show the dilution curve for Mg/Mg vs. Mg/Mg, using the same AGB models as in Ventura et al. (2014) . The amount of AGB material increases from left to right in the plot, with the leftmost point giving the pristine composition assumed in their models. Clearly, the bulk of our measured isotopic ratios fall well below the predicted value, even when considering a composition of purely pristine material. However, if one started from a composition with a lower abundance of heavy Mg isotopes, the curve would shift to the left, and show much better agreement with the observations. We note here that the pristine composition assumed by Ventura et al. (2014) starts out with a higher initial fraction of the heavy Mg isotopes, which would explain this offset. During the burning chain that creates the Al enhancement in AGB stars, production of primarily Mg, but also Mg is occurring, and one would thus expect a change in the isotopic distribution between the pristine and polluted populations. Indeed, the AGB models of Ventura et al. (2014) predict a strong decrease of Mg whereas Mg should increase by almost an order of magnitude. Thus, one would expect the Mg abundances to correlate with [Al/Fe] for the polluted population. As discussed earlier, our analysis does not give any indication of such a correlation being present. This is further illustrated in Fig. 20 , where we plot the dilution curve (solid line) for the Mg isotopic fraction vs. [Al/Fe], using the models of Ventura et al. (2014) . The maximum allowed fraction of AGB material is 0.3, as imposed by the constraints from the He abundance variation measured from the extent of the Horizontal Branch in the cluster. Whereas our [Al/Fe] values are within the model prediction, we see no evidence for a strong increase in Mg. Thus, our observations are somewhat in disagreement with the
Paper: 2018ApJ...866..140Y (fulltext)
sources due to a bias from the FIRST sensitivity. The rest-frame g -band absolute magnitude was estimated by the SED template fitting Mizuki (Tanaka 2015 ; Tanaka et al. 2018 ). There may be two sequences of the g magnitude along the redshifts. The bright sequence with M g ≲ −21 does not strongly depend on the redshift. On the other hand, the number of sources belonging in the faint sequence seems to dramatically increase with increasing redshifts. The sources with faint absolute magnitudes are also faint in apparent magnitudes ( i > 24), and the optical SEDs of the HSC-SSP bands show that some sources on the faint sequence have a flat SED without distinct features. Thus, the redshifts of these sources are probably estimated incorrectly. In fact, the comparison of M g and redshifts with the COSMOS2015 catalog (Laigle et al. 2016 ) in UD-COSMOS shows the underestimated redshifts of the faint sequence sources, and the redshifts of COSMOS2015 indicate that most sources in the faint sequence should be at z L16 > 1.5. In this comparison, a source at z = 0.5 on the faint sequence was found to have redshift and M g consistent with the COSMOS2015 catalog. This source is really a local source with a faint optical luminosity. The faint sequence could be an apparent one due to both the effects of the underestimated redshifts and the Malmquist bias of the HSC-SSP sensitivity limits. Some sources at z ≲ 0.5 may be local galaxies with strong radio radiation from active star formation. Figure 13. Rest-frame 1.4 GHz radio luminosity (upper panel) and the absolute g -band magnitude at the rest frame (lower panel) as a function of photo- z . Only RG sources with reliable photometric redshifts, with z photo > 0, , and , are shown. The black circles and red circles denote the Wide and UD-COSMOS RGs, respectively. Radio loudness at the rest frame was estimated and is defined here as a ratio of the rest-frame 1.4 GHz flux to the rest-frame g -band flux. The wavelength of the rest-frame g -band flux is close to the rest-frame B -band, which is typically used for calculating the radio loudness. The rest-frame radio loudnesses of the SDSS- and HSC-level sources as a function of the apparent i -band magnitude are shown in Figure 14 . We found that RGs with high radio loudness () are found in the HSC-level subsample. Certainly, there are likely also low radio-loudness RGs in the HSC-level regime, but such sources are not included due to the sensitivity limit of the FIRST survey. Note that the radio loudness we use here is underestimated, because HSC band flux could be dominated by stellar light from host galaxies of RGs, and the rest-frame g -band flux is derived from the SED templates of galaxies. Figure 14. Radio loudness at the rest frame as a function of the apparent i -band magnitude. The RG sources with reliable photo- z in the Wide (black) and the UD-COSMOS (red) layers are shown. The blue dotted line is the border separating the SDSS-level and HSC-level samples
Paper: 2016A&A...586A.146V (fulltext)
binary components are not hidden in the glare of an accretion disk. Hence, the short-period hot subdwarf binaries represent an important progenitor channel of CVs. In fact, Schenker (2005) postulates that the CVs below the period gap are the product of post-subdwarf binary evolution. On the other hand, the hot subdwarfs with massive WD companions are potential progenitors of Type Ia supernovae (see Iben Tutukov 1984 ; Webbink 1984 , and for the most recent candidates see Geier et al. 2013 ; Vennes et al. 2012 ). Therefore, they are an important ingredient in understanding Type Ia supernovae. Establishing the accurate masses of both components of post-common-envelope binaries (PCEBs) is the only empirical way to constrain the efficiency parameters of the CEE phase. Furthermore, detailed investigation of hot subdwarf binaries is crucial when determining their masses for comparison with their theoretically proposed evolutionary channels, as well as to test the binary-population synthesis models. However, most of the hot subdwarf binary systems with known orbital periods are single-lined spectroscopic binaries, making it impossible to reliably determine the absolute masses of their components. Among the PCEBs, 15 eclipsing binary systems have been found to date, which are comprised of a hot subdwarf primary (sdO or sdB star) and a low mass secondary, either a late dM or a brown dwarf (BD) (for the most recent discoveries see Schaffenroth et al. 2013 ; Barlow et al. 2013 ; Schaffenroth et al. 2014 , 2015 ). Such eclipsing systems are known as HW-Vir-type binaries after the prototype. In addition to eclipses, HW-Vir systems are notorious for having a very strong “reflection effect”. This reflection effect, characteristic for all close binaries that comprise a very hot primary and much cooler secondary, is due to the high contrast in the temperatures between the heated and unheated hemisphere of the secondary star. As the secondaries in these systems are supposed to be orbiting synchronously, one side of the cool secondary star is constantly being irradiated by the strong ultraviolet (UV) flux of the hot-subdwarf primary. The light-curve solutions of HW-Vir-type systems imply that the temperature of the cool companion around the substellar point is increased up to 10 00020 000 K ( Hilditch et al. 2003 ; Kiss et al. 2000 ). However, the detection of spectral features from the irradiated secondary in HW-Vir-type systems is hampered by the contamination from the much hotter subdwarf primary. Recently, the eclipsing PCEBs gained even more attention, because nearly all systems with accurate eclipse-timing measurements, which have long enough data sets, are found to show apparent orbital-period changes ( Kilkenny 2014 ) that might indicate the presence of circumbinary planets (see Beuermann et al. 2010 , 2012a , b , for the most convincing cases). The claims of circumbinary planets in short-period sdB+dM/BD systems, as well as various planets in close
Paper: 2018A&A...616A..94V (fulltext)
), Since Δ ν is expected to scale as (e.g., Ulrich 1986 ; Brown Gilliland 1994 ) we can substitute Δ ν in Eq. (7) . Assuming Γ 0 ( ν max ) ∝ M c T eff d we can then write, From this development, it can clearly be seen that the b exponent, which represents the Δ ν dependence (Eq. (7) ), does not produce an important influence in the departure from the L / M power law. On the contrary, the influence of the c exponent, introduced by Γ, amplifies the departure from a power law in L/M . In order to simplify the approach, we consider that the relative variation of the effective temperature on the RGB and in the clump is so small compared to the variation in luminosity that we can neglect the T eff dependence. Therefore, we performed a fit of A 0,bol as power law of L and M only. We find a variation in L 0.77±0.06 M −1.25±0.14 with the FREQ method, and in L 0.70±0.06 M −1.40±0.13 with the PROB method. In order to evaluate the validity of the previous hypothesis leading to neglect the T eff influence, the variation of the mode amplitude with L , M and T eff was also measured. We find a variation in L 0.70±0.07 M −1.03±0.15 T eff −4.79±1.36 with the FREQ method, and in L 0.61±0.06 M −1.14±0.14 T eff −5.48±1.18 with the PROB method. The dependence in L and M are less important, as expected, but consistent with the ones we previously found, considering the error bars. In contrast to these results, we can see that the measured T eff dependence is very important. However, the large error bars on the exponent and the lack of important impact on the other dependencies lead us to confirm the simplification hypothesis and neglect the T eff dependence. To assess the impact of the clump stars on the fits, we also consider the RGB stars alone and find significant differences: A 0,bol varying as L 0.84±0.08 M −1.34±0.27 with the FREQ method, and in L 0.82±0.07 M −1.61±0.27 with the PROB method. As stated above, seismic and stellar parameters of stars do not evolve much during the helium burning phase, on the contrary to the RGB phase. It follows that clump stars with similar masses will have similar A 0,bol values, thus making it difficult to disentangle the luminosity and mass dependence for these stars. The fit performed only on RGB stars allows us to get rid of a possible influence of one dependence on the other, and is therefore the one considered throughout the rest of the article. The exponents we find for the luminosity are slightly larger than the ones theoretically predicted ( Samadi et al. 2012 ), but agree with previous measurements ( Mosser et al. 2010 , 2012a ; Huber et al. 2011 ; Stello et al. 2011 ; Corsaro et al. 2013 ; Kallinger et al. 2014 ). A possible explanation for the difference between the predicted and observed exponents is the strong influence of non-adiabatic effects on the mode driving in red giants. Indeed, the theoretical predictions previously mentioned do not rely on fully non-adiabatic models. This type of modeling has not yet
Paper: 2015ApJ...798..120B (acknowledgments)
The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.
Paper: 2015ApJ...810...51M (acknowledgments)
We are grateful for stimulating and helpful discussions with Piero Madau, Monica Colpi, Mitch Begelman, Francesco Haardt, Ralph Pudritz, Chris McKee, Mike Fall, Elena M. Rossi, Melvyn Davies, Milos Milosavljevic, Kevin Schawinski, Luciano Rezzolla, Pedro Montero, Richard Larson, and Christian Reisswig. L.M., D.F., and S.B. acknowledge the hospitality of KITP, where crucial phases of this work were conceived and carried out in two different KITP programs in 2013 and 2014 (“A universe of Black Holes” and “Feedback: Gravity’s Loyal Opposition”).
Paper: 2015A&A...584A.100Z (acknowledgments)
J.P.U.F. acknowledges support from the ERC-StG grant EGGS-278202. The Dark Cosmology Centre is funded by the DNRF.
Paper: 2018MNRAS.475.2978F (fulltext)
see that the optical analyses have larger scatter by a factor of ≳2 in T eff , ≳4 in log g , and ≳4 in [Fe/H]. The larger scatter suggests that our simplified derivation of statistical uncertainties is sufficient until we can calibrate stellar parameters for a larger set of more homogeneous benchmark stars. Well-calibrated stellar parameters would require a more thorough probabilistic model of both StarNet training and StarNet predictions. Therefore, we consider this very good agreement and do not investigate these intrinsic dispersions further at this time. A comparison of Figs 10 and 13 reveals that, when StarNet is trained on synthetic spectra and compared to DR12, the scatter is a factor of ∼2 lower than the residuals between the DR13 and optical analyses. However, due to the different input physics, analysis techniques, spectral resolution, etc. between StarNet, APOGEE, and optical analyses, we cannot directly compare the precision of the stellar parameters returned by StarNet and an optical analysis at this time. Rather, this similar degree of scatter in the residuals indicates that the systematic errors intrinsic to StarNet are comparable to the systematic errors expected from an optical analysis. This indicates that the StarNet model is sufficient for a single survey analysis and can provide a new and exciting tool for the efficient analyses of spectroscopic surveys. 5.2 M dwarfs in DR13 While APOGEE has primarily observed red giant and sub-giant stars, DR13 includes spectra for known M dwarfs. These stars are accompanied with a specific Target Flag: ‘APOGEE MDWARF’. Initially, we had not removed these stars from the data set and found obvious discrepancies between the StarNet results and those provided by ASPCAP DR13. On closer inspection, we found 5900 stars that resembled the known M dwarfs, but that had not been flagged as M dwarfs. Those stars were removed from our data sets, using one of the following cuts: Metallicity parameters of −0.9 (ASPCAP [Fe/H]) 1.0, and where (StarNet [Fe/H] − ASPCAP [Fe/H]) −0.4, or Gravity parameters of (ASPCAP log g ) >3.6, and where (StarNet log g − ASPCAP log g ) −0.3 The ASSET grid provided with the APOGEE DR12 data release does not include synthetic spectra for M dwarfs, and therefore, StarNet cannot be trained on those stars. Furthermore, StarNet cannot be trained on the observed spectra at this time since there are too few M dwarfs in the APOGEE data set (1000 flagged). Our method for flagging M dwarfs is not robust, e.g. some RGB stars may have been removed, while other M dwarfs were not. 5.3 NN considerations One limitation of a deep NN is the necessity of large training sets spanning a wide range of the parameter space. The distribution of the stellar parameters in the training set for StarNet (when trained on observed spectra) are shown in Fig. 14 . From these plots, it is evident that there are fewer spectra for stars at high temperatures, low metallicities, and low gravities. As seen in
Paper: 2016ApJ...826..226K (fulltext)
speed projected onto the beam direction. A.2. Plasma Instability A multipeaked electron distribution function , such as that shown in Figure 2 with the addition of the thermal electron peak at small velocity (not in the figure, that only shows the scattering source) is unstable to the electrostatic “bump-on-tail” instability. This instability is essentially inverse Landau damping, and occurs whenever the derivative of the distribution function with respect to some component of electron velocity is positive. Electron plasma waves with phase velocities ( ω and k are the plasma wave frequency and wave vector, related by the dispersion relation , where is the plasma frequency and the Debye length) for which grow at a rate of (Chen 1974 )where is the normalized one-dimensional electron distribution function of a density n beam of high velocity (Compton recoil) electrons and n e is the background electron density. If the fast electrons are only a small fraction of the total electron density and their speeds are much greater than the electron thermal velocity, conditions that are generally met, then for waves excited by the fast electrons and , where v is a velocity on the rising part of the electron distribution function. Then the growth rate becomeswhere is not normalized to n beam , but is the full electron distribution function. The rapid decrease of drag deceleration with increasing speed, shown by the factor in Equation ( 27 ), acts to create a minimum in (the large peak in from thermal electrons at very small velocity is not shown in Figure 2 ) even when S ( v ) is a monotonically decreasing function, as it can be for power-law photon spectra or a mixture of two- and three-photon annihilation (not shown). For example, if , making the non-relativistic approximation , then . For all but the most steeply decreasing S ( v ), and for v exceeding a few thermal velocities, and the plasma is unstable. In the opposite limit of scattering by very energetic gamma-rays , , and again the plasma is unstable (for quantitative growth rates Equations ( 28 ) and ( 29 ) need to be corrected for the relativistic increase in electron mass). This mechanism of producing unstable electron distributions may operate anywhere an intense flux of gamma-rays, not necessarily from positron annihilation, whose spectrum is not too steep, is incident upon a non-relativistic plasma. The required intensity is determined by the competition between the growth rate γ (Equations ( 28 ) and ( 29 )), proportional to the gamma-ray flux, and the collisional damping rate of the plasma wave that depends on the plasma density and temperature. For a plasma frequency comparable to the frequencies MHz of observed FRBs, is sufficient to grow the instability by ∼10 e -folds to saturation in the sub-millisecond timescale of FRBs. A.3. Energetic Failure This apparently attractive model fails on energetic grounds. The total Klein–Nishina cross-section for a 511 keV photon is cm 2 , so that
Paper: 2017ApJ...841...49C (acknowledgments)
This research was supported by the Korea Astronomy and Space Science Institute under the RD program, Development of a Solar Coronagraph on International Space Station (Project No. 2017-1-851-00), supervised by the Ministry of Science, ICT, and Future Planning. R.-Y.K. acknowledges the support from NASA NNX13AG54G. R.-Y.K. is also supported by the Office of Naval Research and George Mason University.