adsabs/WIESP2022-NER huggingface dataset:
Paper: 2019ApJ...881...56S (fulltext)
however, is insufficient for the pressure gradient to change sign. The top panel of Figure 6 shows the mdot1e-7 case with parameters as in Table 2 , but N 2 = 3 and N 3 = 170. It shows that the active zone inside the rounded-off rim is in its final size. After the active zone has reached a certain size small temperature deviations start to occur and travel through it. These small deviations manifest in the depiction as walls where the dust is present. Since the simulation does not include latent heat it is not possible to decide whether these small deviations have to occur or not, but they were consistently found for every parameter set that produced a dust-free zone. The surface density in Figure 7 mirrors the cavity that is also present in Figure 6 . The dust density falls to g cm −2 in the dust-free zone. To further analyze the equilibrium in the dust-free zone an alternative sublimation formula is in progress. Figure 7. Comparison of gas (solid line) and dust (dotted line) surface density for model mdot1e-7 . The optical depth at constant radiusfor near-infrared radiation in Figure 8 suggests that the dust-free zone forms after a large enough depth in the θ direction is reached because diffusion of radiation energy is no longer sufficient to cool the disk. Figure 8. Vertical optical depth in NIR for model mdot1e-7 . 5. Discussion This work represents an early stage in developing a self-consistent model of the inner rim of protoplanetary disks. In this section some important limitations of the model will be discussed. The simplifications introduced in Flock et al. ( 2016 ), such as a uniform dust-to-gas ratio in the shadowed region and a non-frequency-dependent gas opacity are also applied in this work. A more elaborate sublimation formula is needed in the future, especially with respect to the dust-free cavity inside the disk. A detailed implementation of gas line radiation transfer would go beyond the scope of this work. Further, the value of α was chosen to be constant throughout the box. However, the strength of the turbulence is expected to be influenced by the temperature and could be lower in MRI inactive regions (Lesur et al. 2014 ; Turner et al. 2014 ; Simon et al. 2015 ). A temperature-dependent α will be explored in further work. The temperature deviations and dust walls described at the end of Section 4.3 will have to be further scrutinized in combination with the modeling of the dust-to-gas ratio. Small differences in the temperature can lead to significant changes in the dust-to-gas ratio around the sublimation temperature. This is therefore a delicate problem that needs to be explored tentatively and might necessitate the inclusion of further effects like dust diffusion, latent heat, and temperature-dependent viscosity. 6. Conclusions and Outlook This paper presents a 2D model for protoplanetary disks that consistently includes viscous heating and thermal conduction. The model expands previous works (Flock et al. 2016 )
Paper: 2020ApJ...904..115L (fulltext)
compared to the unambiguously identified [Si ii ] 34.81 and [S iii ] 33.48 μ m makes the [O iv ] identification much more plausible. If [Fe II] 25.98 μ m is not detected, the putative identifications of the intrinsically weaker [Fe II] 35.35 and [Fe II] 17.94 μ m (Smith et al. 2009 ) in the PR equilibrium ejecta are also doubtful. Isensee et al. ( 2012 ) clearly saw the 17.94 μ m line, and we interpret this as coming from the reverse shock. Gerardy Fesen ( 2001 ) saw even higher excitation [Fe II] lines with wavelengths in the range ∼1–2 μ m that are clearly from fast-moving knots or quasi-stationary flocculi and therefore must result from a shock–clump interaction. Koo et al. ( 2016 ) and Lee et al. ( 2017 ) presented mid- and near-IR observations of [Fe II] from ejecta and CSM knots in Cas A. Koo et al. ( 2016 ) demonstrated that in radiative shock models of this emission, most of the [Fe II] radiation comes from the postshock cooling zone, similar to our Figure 5 . Lee et al. ( 2017 ) found ejecta knots with abundances commensurate with O burning, with elements like Si and S present, as well as a sample of ejecta knots showing strong [Fe II] lines but no Si or S. They interpreted these knots as products of complete Si burning leaving almost “pure” Fe, as found in X-rays by Hwang Laming ( 2009 , 2012 ), or in some cases the products of α -rich freeze-out. It is unclear why such ejecta should be in knots, as opposed to diffuse clouds inflated by radioactivity, but compared to SN 1987A, Cas A was a smaller progenitor at explosion, with a presumably faster expansion and lower opacities leading to a reduced effect of this sort. Koo et al. ( 2018 ) estimated a mass of [Fe II]-emitting shocked dense ejecta of ∼3 × 10 −5 M ⊙ . Extrapolating this clumping to the whole inner ejecta is difficult to do with any accuracy, but assuming that the Fe clump density is the same throughout the inner ejecta as it is at the reverse shock, the total Fe mass should be 3 × 10 −5 M ⊙ multiplied by the ratio of the volumes represented by the inner ejecta (6 × 10 56 cm 3 ) and the reverse shock emission. This last piece is ambiguous, but taking it to be cm 3 , where R r = 1.7 pc = 5 × 10 18 cm is the reverse shock radius and the knot diameter 2 r knot ∼ 2.5 × 10 17 cm gives a ratio of volumes of 7.5 and an implied Fe mass in knots of 2 × 10 −4 M ⊙ . This could be uncertain by at least an order of magnitude either way but is likely to be smaller than the Fe mass in the diffuse Fe–Co–Ni bubbles, estimated below. Features coinciding with [Fe V] 20.85 and 36.34 μ m, which would imply Fe masses of order ∼10 −3 M ⊙ , are also visible in the inner ejecta region. These only appear at marginal significance and must wait for better data for confirmation. The [O iv ] 25.89 and [Si ii ] 34.81 μ m morphologies (Smith et al. 2009 ) and the mass estimate of Arias et al. ( 2018 ) suggest that most of the inner ejecta mass expected by Hwang Laming ( 2012 ) are clumped into a volume
Paper: 2017ApJ...851...63S (acknowledgments)
We acknowledge support from the NSF through CAREER award 0847467, and we thank Genevieve J. Graves for making the code LICK_EW available at http://w.astro.berkeley.edu/graves/ez_ages.html . Based on observations obtained at the Gemini Observatory (part of Gemini program GN-2008A-Q-55, processed using the Gemini IRAF package), which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the National Research Council (Canada), CONICYT (Chile), Ministerio de Ciencia, Tecnología e Innovación Productiva (Argentina), and Ministério da Ciência, Tecnologia e Inovação (Brazil).
Paper: 2016A&A...587A.137M (acknowledgments)
The authors thank the anonymous referee for constructive comments. This work was supported in part by the Deutsche Forschungsgemeinschaft (DFG) via SFB 956. G. Busch is member of the Bonn-Cologne Graduate School of Physics and Astronomy (BCGS). L. Moser and S. Smajić are members of the International Max Planck Research School (IMPRS) for Astronomy and Astrophysics Bonn/Cologne. J. Scharwächter acknowledges the European Research Council for the Advanced Grant Program Number 267399-Momentum. M. Valencia-S. received funding from the European Union Seventh Framework Programme (FP7/2007-2013) under grant agreement No. 312789. The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica. This research has made use of the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. This research is also based on observations made with the NASA/ESA Hubble Space Telescope, obtained from the data archive at the Space Telescope Science Institute. STScI is operated by the Association of Universities for Research in Astronomy, Inc. under NASA contract NAS 5-26555.
Paper: 2018A&A...618A...9L (acknowledgments)
KL and IGH are funded via a University of Glasgow Lord Kelvin-Adam Smith Fellowship. IGH is supported by a Royal Society University Research Fellowship. The authors thank the International Space Science Institute (ISSI) for support for the team “Improving the Reliability of Solar Eruption Predictions to Facilitate the Determination of Targets-of-Opportunity for Instruments With a Limited Field-of-View”, where this work benefited from productive discussions. The authors would like to thank N. Gyenge for very helpful discussions about active longitudes and providing the data from Gyenge et al. (2016) .
Paper: 2016MNRAS.463.2904S (acknowledgments)
Support was provided by the National Aeronautics and Space Administration (NASA) through HST grant GO-13787 from the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS5-26555. Support was also provided by the National Science Foundation (NSF) through grants AST-1210599 and AST-1312221 to the University of Arizona. This study used data products from the Southern H-Alpha Sky Survey Atlas (SHASSA), which is supported by the National Science Foundation.
Paper: 2018MNRAS.474.5393B (fulltext)
shown separately (solid lines) as well as the dark matter at z = 8, 9, and10. The bottom panels show the associated scale-dependent bias, b g (defined as the square root of the ratio of the galaxy over dark matter correlation). The errors bars are delete-1 jack-knife errors with each jack-knife subsample corresponding to a subvolume obtained by removing one of the octants from the full volume of BLUETIDES . The one-halo term dominates at scales r 0.12 Mpc h −1 above which the two-halo term takes over. As expected, the brighter SAMPLE-H16 galaxies are more strongly clustered; furthermore, also note that clustering is steeper for the more luminous SAMPLE-H16 galaxies; therefore, as we go to smaller scales the clustering becomes more and more sensitive to the flux limit, with the one-halo clustering showing a significantly steeper dependence on the limiting magnitude as compared to two-halo clustering. These features have been found in observational studies of clustering dependence on luminosity (Kashikawa et al. 2006 ; Savoy et al. 2011 ; Harikane et al. 2016 ). Figure 5. Spatial clustering of galaxies predicted by BLUETIDES at redshifts 8, 9, and 10 from left to right for galaxy samples of different limiting magnitudes m UV corresponding to Bouwens et al. ( 2015 ) B15 and Harikane et al. ( 2016 ) H16 in red and blue lines, respectively. Top panels: Solid lines represent the one-halo and two-halo contributions of the two-point galaxy correlation function. The coloured dashed lines are the total correlation functions. Thin black dashed lines are the dark matter correlation functions. Bottom panels: Scale-dependent bias at redshifts 8, 9, and 10. The dashed lines are the large-scale linear biases. A significant scale dependence in the galaxy bias is seen at scales 0.5 Mpc h −1 ≲ r ≲ 10 Mpc h −1 (well within the two-halo regime) indicating the presence of non-linear effects at scales larger than the length scales associated with individual host haloes. Evidence of scale dependence in the halo bias has been previously reported using N -body simulations (Reed et al. 2009 ; Jose et al. 2016 ). Jose et al. ( 2016 ) (hereafter J16 ) uses simulations with volumes comparable or larger than BLUETIDES to study non-linear bias at redshifts z ∼ 2–5. J16 reports some universal properties of (redshift independent) in the scale-dependent bias and provides a fitting function. In our upcoming paper (Bhowmick et. al. in preparation), we investigate whether their results agree with the predictions of BLUETIDES by making a precise comparison for a wide range of halo mass bins. We find that our predictions are in fairly good agreement (only slightly lower by an amount ≲ 0.3 dex in general) with J16 , extending the validity of J16 results to regimes well beyond z ∼ 2–5; we have discussed this in more detail in Bhowmick et. al. (in preparation). The foregoing emphasize the importance of measuring clustering directly with simulations as (we shall see next) the
Paper: 2019A&A...623A.168S (acknowledgments)
FS, EV, and MR acknowledge the support from the DFG Cluster of Excellence “Origin and Structure of the Universe.” The simulations have been carried out on the computing facilities of the Computational Center for Particle and Astrophysics (C2PAP). SC acknowledges support from Premiale INAF “MITIC” and has been supported by INFN (Iniziativa specifica TAsP). The authors thank Francesco R. Ferraro for the very useful discussion on BSS and for providing advanced information of the color and magnitude distribution in the bulge cluster. Support for MZ and DM is provided by the BASAL CATA Center for Astrophysics and Associated Technologies through grant PFB-06, and the Ministry for the Economy, Development, and Tourism’s Programa Iniciativa Científica Milenio through grant IC120009, awarded to the Millenium Institute of Astrophysics (MAS). EV and MZ acknowledge support from FONDECYT Regular 1150345. DM acknowledges support from FONDECYT Regular 117012. SH, SC, and ES acknowledge grant 309290 form IAC and grant AYA2013-42781P from the Ministry of Economy and Competitiveness of Spain.
Paper: 2019ApJ...871...63M (fulltext)
dwarf—M dwarf eclipsing binaries. We drew sample 1 from Henry McCarthy ( 1993 ), Delfosse et al. ( 2000 ), and Benedict et al. ( 2016 ), excluding those in our calibration sample (Section 2 ). We added one target, GJ 2005A, for which the BC pair was analyzed in this paper, but Seifahrt et al. ( 2008 ) provide a mass determination for the A component. For sample 2, we drew eclipsing binaries from the compilations of Hartman et al. ( 2011 ) and Parsons et al. ( 2018 ), restricting the sample to double-lined eclipsing binaries, those with individual mass estimates better than 5%, and systems with a parallax from the second Gaia data release passing the Gaia quality cuts given in Appendix C of Lindegren et al. ( 2018 ). We removed 19e-3-08413AB because the distance is too large (≫500 pc) to assume zero reddening and excluded systems from Kraus et al. ( 2011 ) because they have no reported flux or luminosity ratios, which are needed to estimate Δ K S (detailed below). Lastly, we removed PTFEB132+19AB (Kraus et al. 2017 ) because it is young (≃650 Myr). In total this gives us individual masses for 29 stars with which to test our – M * relation. A significant advantage of eclipsing binaries is that we could adopt the much more precise parallaxes from Gaia , as they all have orbits that are too tight to show detectable centroid motion. For the astrometric binaries/triples, we drew parallaxes either from Gaia DR2 parallaxes of nearby companions (e.g., for Gl 644ABC we used the Gaia parallax from the wider companion Gl 643) or from sources that accounted for the high-order nature of the system. As with binaries analyzed in this paper, unresolved K S values were taken from 2MASS. We used AO data from VLT/NaCo to derive Δ K for Gl 866AC-B and our own Keck/NIRC2 measurements for Gl 644A-BC and GJ 2005A-B-C (all three are resolved), just as was done for binaries analyzed in this paper. For other systems, including unresolved components of the triples (Gl 866A-C and Gl 644B-C), the literature only provided contrast ratios in optical bands. We converted these to Δ K S using the synthetic and observed magnitudes given in Mann et al. ( 2015 ), following a procedure analogous to that outlined in Kraus et al. ( 2017 ). To briefly summarize, we found the combination of two single-star spectral templates that reproduced both the unresolved spectral energy distribution ( Gaia and 2MASS photometry) and the measured contrast(s) from the literature reference (usually V , R , or Kepler ) and then computed a Δ K S value for the best-fit template combination. This is similar to our conversion of Δ K X to Δ K S detailed in Appendix A . Correction from optical to NIR contrasts depends on metallicity (e.g., Schlaufman Laughlin 2010 ), and precise metallicities of these systems were not known. Instead, we assumed that all systems were −0.6 [Fe/H] +0.4, and we adopted uncertainties that encompass the range of values due to unknown metallicity. Errors introduced from these
Paper: 2015MNRAS.449.1441K (acknowledgments)
We thank the anonymous referee for useful discussions and comments which considerably helped us to improve the manuscript. DK was supported by RFBR grant 13-02-92119 and by the French national research agency ANR (COSMIS project). She thanks the Max-Planck-Institut für Astrophysik for the warm hospitality. SB was supported by RFBR grant 13-02-92119 and SAI MSU. This work was supported by the Deutsche Forschungsgemeinschaft via the Transregional Collaborative Research Center TRR 33 ‘The Dark Universe’ and the Excellence Cluster EXC153 ‘Origin and Structure of the Universe’.
Paper: 2019MNRAS.488.2629A (acknowledgments)
ACKNOWLEDGEMENTS We thank the referee for the constructive comments. Numerical computations were carried out on the Cray XC30 and XC50 at the Center for Computational Astrophysics, National Astronomical Observatory of Japan, and the octopus at the Cybermedia Center, Osaka University. This work is supported in part by the MEXT/JSPS KAKENHI Grant Number JP17H04827 (HY), 18H04570 (HY), and JP17H01111 (KN). YL acknowledges support from National Science Foundation through grants AST-1412719 and MRI-1626251. We are grateful to Dr. Ouchi for helpful comments on the dust temperature. KN thanks to Dr. Kawamata for useful discussions on galaxy size evolution. KN acknowledges the travel support from the Kavli IPMU, World Premier Research Center Initiative (WPI), where part of this work was conducted.
Paper: 2018A&A...615A..68R (acknowledgments)
Acknowledgements This paper is based on observations obtained with MegaPrime/MegaCam, a joint project of CFHT and CEA/DAPNIA, at the Canada-France-Hawaii Telescope (CFHT) which is operated by the National Research Council (NRC) of Canada, the Institut National des Sciences de l’Univers of the Centre National de la Recherche Scientifique (CNRS) of France, and the University of Hawaii. It is based on observations obtained at the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the Science and Technology Facilities Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), Ministério da Ciência, Tecnologia e Inovacão (Brazil) and Ministerio de Ciencia, Tecnología e Innovación Productiva (Argentina). Observations were also obtained with FORS1 and FORS2 at the Very Large Telescope on Cerro Paranal, operated by the European Southern Observatory, Chile (ESO Large Programs 171.A-0486 and 176.A-0589). 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. Funding for SDSS-III has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, and the U.S. Department of Energy Office of Science. The SDSS-III web site is http://www.sdss3.org/ . SDSS-III is managed by the Astrophysical Research Consortium for the Participating Institutions of the SDSS-III Collaboration including the University of Arizona, the Brazilian Participation Group, Brookhaven National Laboratory, Carnegie Mellon University, University of Florida, the French Participation Group, the German Participation Group, Harvard University, the Instituto de Astrofisica de Canarias, the Michigan State/Notre Dame/JINA Participation Group, Johns Hopkins University, Lawrence Berkeley National Laboratory, Max Planck Institute for Astrophysics, Max Planck Institute for Extraterrestrial Physics, New Mexico State University, New York University, Ohio State University, Pennsylvania State University, University of Portsmouth, Princeton University, the Spanish Participation Group, University of Tokyo, University of Utah, Vanderbilt University, University of Virginia, University of Washington, and Yale University. Parts of this research were conducted by the Australian Research Council Centre of Excellence for All-sky Astrophysics (CAASTRO), through project number CE110001020. We thank Mickaël Rigault, Adam Riess, Dan Scolnic, and David Jones for useful discussions.
Paper: 2019ApJ...882...39H (fulltext)
( 1963 ) with the origin located at the barycenter. There is an implicit perturbing force that induces the following accelerationwhich is implicit in the isotropic variation of the mass-losing body. This has been derived in several previous works (Hadjidemetriou 1963 , 1966a , 1966b ; Verhulst 1975 ). If the mass variation is anisotropic, i.e., either from nonconservative mass exchange and asymmetric ejecta, then the equation of motion will have an explicit perturbing force (e.g., Dosopoulou Kalogera 2016a , 2016b ). The orbital evolution can be described as an osculating orbit, which is the Keplerian orbit the binary would follow if the perturbation ceased immediately. The full set of equations for the time evolution of the osculating orbital elements from the isotropically varying mass are (e.g., Hadjidemetriou 1963 ) where the specific angular momentum is conserved, i.e., the quantity is a constant of the motion. Since there is no net force in the direction perpendicular to the orbital plane, , and the orbital motion thus remains coplanar. The mass variation causes the argument of periapsis to precess at a rate (see Equation 3 (c)). 3. Quadrupole Gravitational Radiation For eccentric Keplerian orbits, the orbital separation r between the binary masses obeysThe components of the mass quadrupole tensor I jk for an eccentric binary are where θ = ν + ω , which simplifies to θ = ν if the precession of the argument of periapsis is taken to be negligible. We can choose the initial θ to begin where ω = 0. For isotropic wind mass loss, the periapsis distance is constant, so . When including gravitational radiation reaction, the periapsis distance will indeed decrease, though one can take on timescales less than the mass-loss timescale when the GW coalescence timescale is much longer than the mass-loss timescale. The third time derivatives of the mass-quadrupole tensor components are If the periapsis precession is taken to be negligible, , then the first two terms of Equation 6 simplify to the canonical PM expressions. The additional terms in Equation 6 contain time derivatives of the masses of the binary, which may also be thought of as time derivatives of the binary’s chirp mass . Gravitational radiation carries away the orbital energy and angular momentum of the binary at the following rates where c is the speed of light and is the reduced (trace-free) quadrupole moment tensor . If we take the orbit to be circular, then the GW luminosity simplifies towhere the positive terms in Equation ( 8 ) correspond to GW energy loss from the orbit and the negative terms correspond to GW energy gain in the orbit. From the strain tensor in the transverse-traceless gauge,where is the luminosity distance to the source, the strain amplitude for a circular binary iswhere Ω is the angular frequency of the binary, . In the transverse-traceless gauge, the plus and cross polarizations of the strain are and . From Equation ( 10 ), we can define the following parameters