adsabs/WIESP2022-NER huggingface dataset:
Paper: 2017ApJ...834..135T (fulltext)
with high stellar mass surface densities (see next section). Figure 5. Comparison of circularized half-light radii at ALMA/870 μ m (red circles) with those at HST / H 160 band (cyan squares) for our sample of 12 galaxies with size measurements at 870 μ m. Gray circles indicate the rest-optical sizes for star-forming galaxies around the main sequence at z = 1.9–2.7, which are drawn from the 3D- HST survey. A dashed line shows their fitting function, . Table 2 Galaxy Properties for 12 SFGs with 870 μ m Size Measurements Notes. a Sérsic indices and half-light radii at 1.6 μ m. We do not use U4-16795 because the best-fit Sérsic index reaches the constrained limit of n = 8. b Half-light radii at 870 μ m. c Stellar mass surface density within a central 1 kpc calculated in stellar mass maps. d SFR surface density within a central 1 kpc calculated from the best-fit exponential models at 870 μ m and total SFRs. e Bulge formation timescales to reach the stellar mass surface density of log(Σ M bulge / M ⊙ kpc −2 ) = 10 (Equation ( 3 )). f Gas depletion timescales by star formation and outflows (Equation ( 4 )). In the analysis of size measurements, we do not include two massive star-forming galaxies with log( M * / M ⊙ ) > 11 in the parent sample of galaxies identified by the narrowband survey. One is not observed with ALMA, and the other one (U4-36568) is not detected in the high-resolution map (Figure 1 ). Given the high completeness of 86% (12/14) in the stellar mass range, our results are not significantly affected by the sample selection. Therefore, we find massive galaxies to commonly form stars in the extremely compact central region as at least 64% (9/14) have small 870 μ m sizes of R 1/2,870 μ m 1.5 kpc. This result is in excellent agreement with similar and independent evidence coming from an ALMA/870 μ m study of six massive star-forming galaxies at z ∼ 2 (Barro et al. 2016 ). Barro et al. ( 2016 ) find that the mean half-SFR radius is ∼30% smaller than the mean half-mass radius. The main difference between our work and Barro et al. ( 2016 ) is that they preselect only optically compact star-forming galaxies while our study almost completely selects main-sequence galaxies. 4. BULGE FORMATION IN EXTENDED, ROTATING DISKS Given that our ALMA sample is already massive, , they are likely to soon thereafter quench the active star formation and to be observed as quiescent galaxies in the local universe. Quiescent galaxies are always smaller than star-forming galaxies at any redshift and any stellar mass and have a cusp profile ( n > 2) unlike star-forming galaxies with exponential disks (Wuyts et al. 2011b ; van der Wel et al. 2014b ). A spatial distribution of stars within galaxies would not be changed unless a violent process like major mergers happens. However, centrally concentrated star formation reduces the half-light or half-stellar-mass radii of galaxies, and their Sérsic index would increase by central bulge formation. We quantitatively
Paper: 2019ApJ...883...78P (acknowledgments)
We would like to thank John O’Meara for providing HIRES spectra and Nigel Mathes for reducing UVES spectra. We thank Nicolas Lehner for discussions on the UV ionizing background. We also thank Neil Crighton for the MCMC analysis software and Cloudy ionization training. Support for this research was provided by NASA through grants HST GO-13398 from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. S.K.P. acknowledges support through the Australian Government Research Training Program Scholarship. G.G.K., N.M.N., and M.T.M. acknowledge the support of the Australian Research Council through the Discovery Project DP170103470. C.W.C. and J.C.C. acknowledge support by the National Science Foundation under grant No. AST-1517816. Parts of this research were supported by the Australian Research Council Centre of Excellence for All Sky Astrophysics in 3 Dimensions (ASTRO 3D), through project number CE170100012. Some of the data presented herein were obtained at the W. M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration. Observations were supported by Swinburne Keck programs 2014A_W178E, 2014B_W018E, 2015_W018E and 2016A_W056E. The Observatory was made possible by the generous financial support of the W. M. Keck Foundation. The authors wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Maunakea has always had within the indigenous Hawaiian community. We are most fortunate to have the opportunity to conduct observations from this mountain. Based on observations collected at the European Organisation for Astronomical Research in the Southern Hemisphere under ESO programs listed in Table 2 .
Paper: 2015MNRAS.454.2691M (acknowledgments)
We thank Tsang-Keung Chan, Xiangcheng Ma, Daniel Anglés-Alcázar, Freeke van de Voort, Robert Feldmann, Chris Hayward, and Michael Anderson for various constructive suggestions. We thank Romeel Davé for providing details on wind implementation for our simulation comparison section. We thank the anonymous referee for providing suggestions that improved this work. DK was supported by a Hellman Fellowship and NSF grant AST-1412153, and by funds from the University of California San Diego. Support for PFH was provided by the Gordon and Betty Moore Foundation through Grant 776 to the Caltech Moore Center for Theoretical Cosmology and Physics, by the Alfred P. Sloan Foundation through Sloan Research Fellowship BR2014-022, and by NSF through grant AST-1411920. C-AF-G was supported by NSF through grant AST-1412836, by NASA through grant NNX15AB22G, and by Northwestern University funds. EQ was supported by NASA ATP grant 12-APT12-0183, a Simons Investigator award from the Simons Foundation, the David and Lucile Packard Foundation, and the Thomas Alison Schneider Chair in Physics at UC Berkeley. The simulations analysed in this paper were run on XSEDE computational resources (allocations TG-AST120025, TG-AST130039, and TG-AST140023). We would like to thank the Simons Foundation and the participants of the Galactic Super Winds symposium for stimulating discussions. We would also like to thank the Kavli Institute for Theoretical Physics and the participants of the Physics of Star Formation Feedback program for interactions that improved this work.
Paper: 2020ApJ...889...33Z (fulltext)
Other channels may account for the most carbon-rich early-R stars. 4. Conclusion and Discussion The merger of an HeWD with an RGB star in a CE provides a possible model for the origin of lithium- and carbon-rich stars. By adopting a grid of approximate initial conditions for such mergers informed by previous SPH simulations, we have made one-dimensional mesa calculations of the post-merger evolution of a wide range of possible systems. Using BSPS we have computed the frequency and distribution of HeWD + RGB binary-star systems that are likely to merge. Analysis of the post-merger evolution models show that the final surface abundances depend on the masses of the progenitor HeWDs. The Li-rich giants form from low-mass HeWDs, i.e., ; the C-rich giants form from low-mass HeWDs, i.e., . The pathway to either of these categories is related to the temperature of the helium-core-burning shell and the balance of nuclear products during initial shell burning. Analysis of the distribution of post-merger systems, as given by BSPS and mesa evolution tracks, shows that predictions for HeWD+RGB mergers are consistent with observations of most Li-rich giants in terms of surface effective temperature ( T eff ), surface gravity (log g ), surface luminosity (log L ), surface abundance ( A (Li) and 12 C/ 13 C), and the Galaxy disk space density of . They are also consistent with the properties of some early-R carbon stars, in terms of T eff , log g , C/O, A (Li), 12 C/ 13 C, and a space density of (3.5−5.9) × 10 −7 pc −3 . According to the models, the post-merger objects are helium-core-burning stars, which means that they are all in the red clump stage of evolution, i.e., they are red giants with nondegenerate helium cores. However, only a few Li-rich giants have had their precise evolutionary stage determined by asteroseismology. Those stars that have been analyzed in this way correspond well with our results. Hence, we strongly argue that the Li-rich red clump stars and the majority of early-R carbon stars most likely form following the merger of an HeWD with an RGB star. A few problems still need to be addressed: (i) Post-merger evolution calculations depend on evolution during the CE phase that is determined by the progenitors of the merger, the amount of mass loss during the CE phase, and the structure of the CE. However, our knowledge of CE phase in binary-star evolution is still poor. Future theoretical studies and simulations will improve this deficiency. Also, as the products of post-CE mergers, further observations of Li-rich red clump stars will provide additional data for studies of the CE phase. (ii) Some of the observed Li-rich giants have measured rotation velocities, but we have not calculated rotation speeds for post-merger models immediately following CE ejection. Our model will be improved with more knowledge of angular momentum transport during the merger. (iii) The HeWD+RGB merger model provides a successful explanation for Li-rich red
Paper: 2020A&A...634A..28L (acknowledgments)
The authors are grateful to François Mignard for his useful comments and suggestions which improve greatly the manuscript. The work of N.L. has been supported by China Scholarship Council (CSC) for the joint doctor training program at Department SYRTE (Systèmes de Référence Temps-Espace) of the Paris Observatory (File No. 201706190125). N.L. is also funded by the National Natural Science Foundation of China (NSFC) under grant No. 11833004. N.L. is grateful to Nan Jiang for her careful check of the language. 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 research has made use of NASA’s Astrophysics Data System. This research also made use of Matplotlib, an open-source 2D graphics package for Python ( Hunter 2007 ) and Astropy ( http://www.astropy.org ), a community-developed core Python package for Astronomy ( Astropy Collaboration 2013 , 2018 ).
Paper: 2018MNRAS.481L..69S (fulltext)
high-redshift probe of the first stars. Figure 2. Confronting structure formation theory with observations. Shown is the critical (minimum) DM halo mass required for the runaway collapse of primordial gas versus redshift for different cosmological models (see Fig. 1 ). The three coloured lines represent the critical halo properties for CDM ( blue ) as well as FDM models with ( orange ) and ( green ). The symbols indicate 2σ, 3σ, and 4σ peaks, respectively. The magenta region summarizes estimates for the DM halo mass, hosting UFD galaxies in the Local Group. Complementary to this, the light-grey region delineates the redshift range inferred from the observed EDGES spectral trough in the 21-cm absorption profile (Bowman et al. 2018 ). The dark-grey region shows the epoch of reionization. As is evident, the combined empirical constraints disfavour FDM models with the lightest axion masses. 3.1 Local Group dwarf galaxies The first constraint is the mass limit for the first star-forming haloes, M crit , such that at a given redshift runaway cooling can only occur in sufficiently massive hosts. Observations of UFD galaxies provide non-trivial tests for any model of structure formation, as the model must allow star formation in low-mass host haloes, sufficiently early in cosmic history. Possibly, masses of the UFD progenitors may be as small as at a redshift of z ≈ 10 (Bland-Hawthorn, Sutherland Webster 2015 ). However, observations of individual UFDs in the present-day Universe give relatively uncertain virial masses of the order of (e.g. Wolf et al. 2010 ; Simon et al. 2011 ; Boylan-Kolchin, Bullock Kaplinghat 2012 ), and a recent statistical analysis of local UFDs provides a conservative upper limit of the minimum UFD-hosting halo mass of (Jethwa, Erkal Belokurov 2018 ). We illustrate the plausible range of UFD halo masses in Fig. 2 . Structure formation models represented by curves that lie beneath the upper limit of this region are consistent with the Local Group constraints. In the ΛCDM cosmology, M crit is always below the region constrained by the UFD observations, regardless of overdensity νσ. This is in accordance with the standard picture of hierarchical structure formation. For the FDM model with the lighter axion, , on the other hand, M crit lies above the allowed UFD region for 2σ peaks. In this case, the most frequently occurring haloes would form with masses larger than what is inferred for the UFD hosts. The resulting tension implies that existing UFD observations disfavour the lighter-axion FDM model. Conversely, this tension does not exist in the model with heavier-axion FDM (). Here, values of M crit for reasonable, not too unlikely, σ peaks lie within or below the allowed UFD region. Though it remains unclear whether UFDs alone can rule out FDM entirely, further restrictions for the allowed models are anticipated through observations of lower-mass UFDs with upcoming large-scale surveys, carried out, e.g., with the Large Synoptic
Paper: 2019MNRAS.490.3882S (fulltext)
derived from CfA3 in red. 5.3 Comparison of light-curve fitter results To make any cosmological statements based on the results in the previous section is beyond the scope of this paper, as this would require a detailed study and justification of the utilized light-curve fitters and their parameters, among many other considerations. It is interesting and possible, however, to compare results from the two light-curve fitters we employ to check for consistency. As the principal quantity of interest when fitting the light curves of SNe Ia is distance, we will focus our comparison on the derived distance moduli. The left-hand plot in Fig. 12 compares the distance moduli from snoopy and mlcs2k2 after correcting to put the measurements on the same scale (so that relative distance moduli are compared, independent from assumptions about the Hubble constant). This correction consists of adding an offset to the distance moduli from each fitter such that the value of H 0 measured from each set of results yields 65 km s −1 Mpc −1 . We perform this comparison only for spectroscopically normal SNe Ia in our sample for which snoopy finds Δ m 15 1.7 mag and for which z > 0.01. Of course, further restrictions should be placed when selecting a sample for cosmological purposes, but our selection is reasonable for performing a general comparison. We find strong agreement between the two sets of corrected distance moduli – a Kolmogorov–Smirnoff test gives a P -value of 1.000. The median residual is −0.026 mag with a statistical dispersion of 0.135 mag. Figure 12. Comparison of the (scaled) distance modulus and host-galaxy reddening results from both light-curve fitters for the selected subset of our data set. If we were to ensure consistency in choosing the parameters for each light-curve fitter, the residuals would almost certainly decrease. In particular, when fitting with mlcs2k2 , we place an exponentially decaying prior on A V , but no such prior was imposed with snoopy . This difference may well manifest in statistically and systematically different results for host-galaxy reddening and distance moduli between the two fitters. We compare host-galaxy reddening results in the right-hand panel of Fig. 12 , where for mlcs2k2 we have converted to reddening using E ( B − V ) host = A V / R V , with R V = 1.7. The agreement is reasonable, with a median residual of −0.056 mag and statistical uncertainty of 0.055 mag. Furthermore, the facts that the median residual ( snoopy minus mlcs2k2 ) is negative and that the disagreement is most severe for small E ( B − V ) host are consistent with what one might expect given the prior imposed by mlcs2k2 . 6 CONCLUSION In this paper, we present BVRI (along with some unfiltered) light curves of 93 SNe Ia derived from images collected by the LOSS follow-up program primarily over the interval from 2009 to 2018, but with several instances as early as 2005. Careful and consistent observational and processing techniques ensure
Paper: 2019MNRAS.487.5106M (fulltext)
gravitational instabilities (GI) operating in the outer disc, driving material towards the protostar, and the magneto-rotational instability (MRI) operating in the inner disc region around the young protostar when the conditions are appropriate for a sufficient ionization fraction (Armitage et al. 2001 ; Zhu et al. 2009a , b , 2010a ; Zhu et al. 2010b ; Stamatellos et al. 2011 ; Mercer Stamatellos 2017 ). As the disc increases in mass by accreting infalling gas from the envelope it becomes gravitationally unstable, as suggested by the formation of the spiral structures seen in the quiescent phase snapshots in the left column of Fig. 1 . Thus, angular momentum is redistributed outwards efficiently resulting in inward gas flow. The inner disc is hotter and gravitationally stable such that there is no mechanism available to redistribute angular momentum, and material accumulates around the protostar (∼1 au). As more gas flows into this region, compressive, viscous, and protostellar heating raise the local temperature further. Once a threshold temperature of 1400 K is reached, this region is sufficiently ionized for the MRI to be activated (Zhu et al. 2009a , b , 2010a ; Zhu et al. 2010b ). The MRI increases the viscosity (the effective viscosity parameter becomes α MRI = 0.1; Shakura Sunyaev 1973 ), providing a way to redistribute angular momentum and resulting in rapid gas accretion onto the protostar. During the outburst phase, the accretion luminosity increases dramatically (∼10 3 L ⊙ ), heating the disc and disrupting the spiral structure (outbursting phase snapshots in the left column of Fig. 1 ; see also MacFarlane Stamatellos 2017 ). Once the gas from the inner region has been accreted, the protostar returns to the quiescent phase. The process of quiescent to outbursting phase repeats multiple times during the evolution of the YSO (Stamatellos et al. 2012 ). Outbursts happen every ∼1000 yr initially, whereas at later stages (but still during the Class O/I stage) every ∼10 000 yr (see Stamatellos et al. 2011 , 2012 ). This is similar to what is inferred from observations (e.g. Scholz, Froebrich Wood 2013 ; Hillenbrand Findeisen 2015 ; Fischer, Safron Megeath 2019 ; see also Peña, Naylor Morrell 2019 , for longer outburst time-scale estimates for older, Class I/II, YSOs). It is assumed that in the quiescent phase gas accretes onto the protostar at a rate of (this is a free parameter in our calculation) and the protostellar luminosity is then on the order of ∼1 L ⊙ . As a comparison, at the same time the accretion rate onto the outer disc is on the order of ∼10 −5 −10 −4 M ⊙ yr −1 . During the outburst phase the accretion onto the protostar is significantly enhanced. This model is different than the models of Vorobyov Basu ( 2005 , 2006 , 2010 ). In their models the GIs alone are responsible for the observed outbursts: GIs create clumps in the disc of a young stellar object that eventually fall onto the protostar, resulting in outbursts
Paper: 2020MNRAS.497..933H (fulltext)
time variations appear for all radii, with typical amplitude changes in η R / η max of ∼0.2, except for the innermost radial bin. The thick solid grey curve indicates the mean bar half-length variations (black curve from Fig. 5 ), showing a very good agreement with the amplitude fluctuations. Note that as the bar decreases in length, so does its amplitude. Very similar time variations appear for all radii, with typical amplitude changes in η R / η max of ∼0.2, except for the innermost radius considered. The thick solid grey curve in Fig. 7 represents the mean bar half-length variations (black curve from Fig. 5 ), which can be seen to agree very well with the amplitude fluctuations, including the short and long periods. The fluctuations seen in the bar amplitude at fixed radii (Fig. 7 ) are more similar to the bar length time variations than the η max , which can vary with radius (see Fig. A6 ). This can generally be seen for Model2 as well in Fig. 8 . One key difference between the two simulations is the fact that for Model1 all radii peak at nearly the same time, while in Model2, they are delayed with the lowest radius of 3 kpc always peaking last. Such a pattern suggests a spiral arm contribution. Indeed, we established in Section 3.4 that the Model2 bar true length is about 3 kpc; therefore, the region examined in Fig. 8 lies at or outside the true bar, yet in the regime where our three measurement methods detect a bar. The solid grey curve in Fig. 7 shows the bar length time variations, seen to follow the overall trend in η R / η max , in best agreement with the two outermost radii. Figure 8. As Fig. 7 , but for Model2. The solid grey thick curve indicates the mean bar half-length variations (black curve from Fig. 6 ), showing an overall good agreement with the amplitude fluctuations. Since the true bar length is ∼3 kpc, the outermost three radial bins for which η R was estimated correspond to the spiral arms, as evident from the systematic time offset among different curves. 4.4 Bar pattern speed We estimated the instantaneous bar pattern speed using the modified TW method by Sanders et al. ( 2019 ), who applied it to both MW data and N -body simulations. The top left-hand panel of Fig. 9 shows the configuration we used to measure Ω p , over the stellar density of Model1. As done by Sanders et al. ( 2019 ), we assumed a bar angle of 33°, a Galactic latitude range | b | 5°, and a solar Galactocentric distance of R 0 = 8.12 kpc. The rotation is clockwise. Estimates are done in bins of Galactic longitude dl = 2° in the radial range indicated by the two arches centred on the Sun (at distances of 4.12 and 8.12 kpc), which is sitting at ( x, y ) = (8.12, 0). The straight black line over the bar has a half-length of 3.2 kpc, which corresponds to l ∼ 17° for our bar angle. The bar angle is kept the same at each time output. The white circles show the bar time-median CR, 4:1 OLR, and 2:1 OLR at R CR = 3.25, R 4: 1OLR = 4.2, and R 2: 1OLR =
Paper: 2017ApJ...836...96M (acknowledgments)
We thank an anonymous referee for the constructive report. L.F.R. and L.L. acknowledge the support of DGAPA, UNAM, and CONACyT, México.