adsabs/WIESP2022-NER huggingface dataset:
Paper: 2020ApJ...890...33B (fulltext)
observations presented in Prochaska et al. ( 2017 ; hereafter P17 ). These data were previously analyzed using PIE models, which, as stated, are usually able to reproduce the low and intermediate ions but have trouble matching O vi (Lehner et al. 2013 ) without invoking very-low densities (Stern et al. 2016 ). The structure of this work is as follows. In Section 2 we discuss the observational data, and in Section 3 we give a brief description of MAIHEM as well as the MCMC code used to conduct our analysis. In Section 4 we show the main results of our analysis, followed by a discussion in Section 5 . 2. Data The data used for this analysis are a compilation of absorption-line systems probing the CGM of low-redshift galaxies from the COS–Halos survey. Specifically, we model the observations presented in P17 , which were previously analyzed using PIE models. These data consist of several ions for multiple absorbers along quasar sightlines that probe M ⊙ galaxies out to an impact parameter b ≈ 150 kpc (Tumlinson et al. 2013 ). The reader is encouraged to refer to previous work by the COS–Halos survey for additional details on the data collection and processing. Here, we study the same COS–Halos sightlines that were analyzed in P17 , for a total of 44 low-redshift ( z ≲ 0.5) systems. These sightlines exhibit absorption from a mixture of ionization states that include C ii , C iii , Fe ii , Fe iii , Mg ii , N ii , O i , O vi , Si ii , Si iii , Si iv , and S iii . As we now have the ability to model a multiphase medium, we also include Si iv and O vi , the column density for which were not incorporated into the analysis of P17 . We apply the MCMC analysis to the entire sample, finding, however, that only 32/44 systems converge on solutions with the MAIHEM model. Systems that are unable to find solutions have either all upper or lower limits, or only have accurate column density measurements for a single ion. When drawing comparisons with P17 , we thus limit our sample to only include the 32 systems with a converged solution. 3. Methods 3.1. MAIHEM Models We use the MAIHEM code to model ionic fractions in an isotropic turbulent medium. This cooling and chemistry package is a modified version of the open-source hydrodynamics code FLASH (Version 4.4) (Fryxell et al. 2000 ), which explicitly tracks the reaction network of 65 ions, including: hydrogen (H i and H ii ), helium (He i –He iii ), carbon (C i –C vi ), nitrogen (N i –N vii ), oxygen (O i –O viii ), neon (Ne i –Ne x ), sodium (Na i –Na iii ), magnesium (Mg i –Mg iv ), silicon (Si i –Si vi ), sulfur (S i –S v ), calcium (Ca i –Ca v ), iron (Fe i –Fe v ), and electrons from an initial non-equilibrium state to a steady state. Modifications for solving the hydrodynamic equations include using an unsplit solver based on Lee ( 2013 ), as well as a hybrid Riemann solver that utilizes the Harten Lax and van Leer solver (Einfeldt et al. 1991 ) to ensure stability in regions of strong shocks, and
Paper: 2017MNRAS.467.2113W (fulltext)
currently say whether IC10 is an example of the radioFIR conspiracy as the sample size used to determine the radioSFR relation is small with a large scatter. Future studies with a larger sample would reveal the true significance of the deviation. 5 CONCLUSIONS We present high-resolution e-MERLIN observations of the nearby dwarf irregular galaxy IC10 as a part of the LeMMINGs project. Our main conclusions are as follows. Our e-MERLIN observations reveal 11 compact radio sources at 1.5GHz coincident with IC10s star-forming disc, five of which also have detections at 5GHz. Although we cannot conclusively classify all sources on the basis of their spectral index due to differences in sampling the uv plane, we gain insight into the nature of each source through the use of low-resolution multiwavelength data, the positional coincidence with features at other wavelengths such as H and 70m, and an estimate of their brightness temperature. Based on our analysis of the expected background source counts, the number of detected sources is consistent with them all being background sources. However, as some sources closely trace compact emission in maps at other wavelengths, giving us confidence that some of these sources do in fact belong to IC10. We find that of the 11 sources, three likely reside in IC10 (sources 1, 2 and 4), all of which are H ii regions; five are background sources (sources 3, 5, 6, 9, 10) and three are yet to be conclusively classified (sources 7, 8, 11). Based on the time-scale over which SNR are typically observable and the resolution limitations of these high-resolution interferometry observations, we expect that we will not be able to detect any SNR with the present e-MERLIN observations. The only SNe we would expect to be able to detect are very young RSNe, but it is unlikely that we would observe any due to a combination of their short observable time-scales and low SFR observed in IC10. We explore other possible factors that could affect SNR detectability in dwarf irregular galaxies and conclude that e-MERLIN alone is well suited to study SNR in galaxies located at distances greater than 4Mpc due to the spatial frequencies covered by the array. For galaxies closer than 4Mpc, additional VLA observations are required to fully cover the spatial scales that SNR candidates occupy. The same reasons as mentioned above explain why we see no compact sources within the NTSB in IC10 and that the detected sources that belong to IC10 are likely to be compact H ii regions. Based on starburst99 models, we determine lower limits for the stellar cluster masses required to produce these observed compact H ii regions; the largest one is about half the size of 30Doradus in the LMC. We present updated integrated flux densities for IC10, where we subtract the contribution from background sources from literature data. We find a 4percent reduction in the measured flux at 1.5GHz and a 2percent reduction at 4.85GHz. With this updated measurement, we
Paper: 2019MNRAS.489.2097B (acknowledgments)
ACKNOWLEDGEMENTS SB thanks Martin Millon and Aymeric Galan to help identifying the issue with astrometry while working on the Time-delay Lens Modelling Challenge (Ding etal. 2018 ). We thank Adriano Agnello, Anowar Shajib, Chih-Fan Chen, Dominique Sluse, Patrick Kelly, Xuheng Ding, and Jenny Wagner for useful feedback on the manuscript. We thank the anonymous referee for useful suggestions and comments that improved the presentation of this work. Support for this work was provided by NASA through grant numbers HST -G0-14254 and HST -G0-15320 from the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS 5-26555. TT acknowledges support by the NSF through grant AST-1714953 and by the Packard Foundation through a Packard Research Fellowship.
Paper: 2018ApJ...868..136M (fulltext)
a spherical explosion triggered inside a rotationally supported torus (presumed to be formed by preceding binary interaction), which results in a collimation of ejecta toward the system’s poles. Frank et al. ( 2018 ), García-Segura et al. ( 2018 ), and Reichardt et al. ( 2018 ) recently discussed the interaction of stellar winds with the debris of a common-envelope episode and how this process might contribute to observed bipolar planetary nebulae. Here we expand on this line of work to show that bipolar morphologies can, instead, arise from the interaction of successive components of equatorial mass loss from the coalescing binary. We analyze numerical simulations of coalescing binaries transitioning from Roche lobe overflow toward the onset of the common-envelope phase. We show that early outflows, contrary to previous theoretical predictions, are bound to the binary system and form a circumbinary torus. Later ejecta, generated just as the secondary object plunges within the envelope of the donor, have larger radial velocities and are unbound. The interaction of unbound ejecta with the bound torus redirects ejecta toward the poles, laying the foundation for the generation of bipolar remnants of stellar coalescence recently reported by Kaminski et al. ( 2018 ). 2. Bipolar Remnants of Stellar Coalescence and Binary Interaction 2.1. Remnants of Galactic Red Novae There is growing evidence that the remnants of galactic luminous red nova outbursts, believed to originate from stellar coalescence, can exhibit bipolar structures. There is strong evidence for a bipolar morphology of V4332 Sgr. We also review the less extensive evidence related to the two other galactic transients that have been studied, V1309 Sco and V838 Mon. Some of the first evidence for bipolar morphologies came from V - and R -band linear polarization measurements of the remnant of V4332 Sgr, which flared in 1994 (Kamiński Tylenda 2011 ). Correcting for the effect of significant unpolarized line emission, Kamiński Tylenda ( 2011 ) estimated continuum polarizations of approximately 40% and 20% for the V and R bands, respectively, and suggested that dust scattering in a bipolar outflow is one of the only morphologies that can produce such high linear polarization in an unresolved source. Spectropolarimetric observations by Kamiński Tylenda ( 2013 ) confirmed Kamiński Tylenda’s ( 2011 ) suggested continuum polarization. Earlier conclusions regarding morphology have recently been strengthened by submillimeter observations, which have offered an opportunity to spatially resolve galactic remnants of stellar-coalescence transients (Kaminski et al. 2018 ). Spatially resolved observations using the ALMA of V4332 Sgr produce a position–velocity diagram that is best modeled with a bipolar molecular outflow with a 60° opening angle (see Figures 5 and 6 of Kaminski et al. 2018 ). The merger of V1309 Sco has been extremely useful in decoding observations of luminous red novae in general.
Paper: 2015MNRAS.447.2631V (acknowledgments)
This research was supported by the grant AYA2012-39303 and SGR2014-1073 (DV, DFT) and the Project Formosa TW2010005, for bilateral research between Taiwan and Spain (KH and DFT). KH is partly supported by the Formosa Programme between National Science Council in Taiwan and Consejo Superior de Investigaciones Cientificas in Spain administered through grant number NSC100-2923-M-007-001-MY3. The research leading to these results has also received funding from the European Research Council under the European Union's Seventh Framework Programme (FP/2007-2013) under ERC grant agreement 306614 (MEP). MEP also acknowledges support from the Young Investigator Programme of the Villum Foundation. We are grateful to the referee for a careful reading and the useful comments.
Paper: 2018MNRAS.481.5098M (acknowledgments)
ACKNOWLEDGEMENTS Based on observations with the NASA/ESA Hubble Space Telescope , obtained at the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS 5-26555. JA and AB acknowledge support from STScI grants GO-13297. This work has been supported by the European Research Council through the ERC-StG 2016, project 716082 ‘GALFOR’ and by the MIUR through the the FARE project R164RM93XW ‘SEMPLICE’. AFM acknowledges support by the Australian Research Council through Discovery Early Career Researcher Award DE160100851.
Paper: 2018ApJ...864...42C (fulltext)
the acceleration of the star induced by the companion. We augment the internal RV uncertainties listed in Table 2 with a “jitter” term that is added in quadrature to accommodate the expected level of noise introduced from stellar variability. HD 169889 A has a visual color of B − V = 0.75 and median activity index of S HK = 0.18. We assume a B − V uncertainty of 0.01 mags. Using empirical relations from Isaacson Fischer ( 2010 ), we estimate a jitter value of σ = 2.2 ± 0.1 m s −1 . Table 2 Relative Stellar RV Measurements for HD 169889 Note. A horizontal line denotes the division between data sets when the HIRES detector was upgraded (2004 August) requiring a relative RV offset for analysis. Table 3 Summary of High-contrast Imaging Observations Showing the Number of Coadds ( n coadds ), Frames ( N fr ), Exposure Time per Frame (Δ t exp ), and Change in Parallactic Angle Δ π (°) Table 4 Summary of Astrometry Measurements Note. Position angle (PA) measurements are referenced with respect to the equator of the observing epoch. Table 5 Photometric Results and Companion Properties A straight line is fit using Markov Chain Monte Carlo (MCMC) methods to find a best-fitting RV slope of m s −1 yr −1 . Combined with the direct imaging projected separation, we use the relation from Liu et al. ( 2002 ) to derive a lower limit to the companion mass of m min = 0.369 ± 0.010 M ⊙ (Table 5 ). Clearly, HD 169889 B cannot be a brown dwarf but is instead a compact object provided the acceleration is attributable to this object. A more careful treatment would consider the slight curvature noticeable by eye in the RV time series as well as astrometry (Figure 1 ). However, only a small fraction of the companion orbit has been traced in either data set. Further, the astrometry measurements were obtained with two different instruments using three different filters. The systematics that result from combining these observations can be noticed in the zoomed-in portion of Figure 5 at the ≈1 σ level (10 mas). The WD companion appears to be moving to the north, i.e., closer to HD 169889 A, but further AO imaging is required to confirm. Figure 5. Proper motion and parallactic motion analysis taking into account the path that an infinitely distant background object would traverse is compared to the separation of HD 169889 B over time. Angular offsets are measured relative to the optical primary star. The inset shows a zoomed-in view of individual astrometric measurements. The companion shares the same space motion demonstrating that they are gravitationally associated. 5. WD Properties Empirical relations between WD masses and their progenitors have been established based on the study of open clusters, globular clusters, and common proper motion pairs (Weidemann 2000 ; Williams et al. 2009 ). Using the dynamical mass lower limit from RVs and imaging, we attempt to place a lower limit on the companion progenitor mass ( M i ) by inverting the empirical relationship from
Paper: 2019MNRAS.486L..52S (fulltext)
time like Fig. 2 , we can conclude that the bending in period B must exist from the observational data about the periatron time T N in both periods A and C. It is possible to judge whether KL-oscillation had happened or not by using our analysis without the observational data in highly eccentric state. One may be worried about the spin evolution of the pulsar because the spin–orbit coupling appears in GR (Barker, O’Connell 1975 ), which may change the direction of the pulsar rotation axis. If the beaming direction of pulse signal is changed, the pulsar will disappear in stage B. However, following Liu Lai ( 2017 ), we find that the adiabaticity, which measures the ratio of the KL oscillation time-scale to the presession time-scale, is very small in the present system. It means that the spin of the inner binary evolves ‘non-adiabatically’ even when the KL-oscillation occurs. As a result, the spin will be parallelly transported just as Newtonian case, and then the beaming direction will not change so much even in the KL-oscillation stage and the whole period (A, B, and C) will be observed. Finally we shall discuss the possible parameter range that the bending of the cumulative-shift curve occurs within our lifetime. In addition to the model given in Table 1 , we have performed our calculation for 19 200 models by changing the outer semimajor axis () and the mass of the third body (10M ⊙ ≦ m 3 ≦90M ⊙ ) and analysed in which parameter ranges the above bending will occur within 100 yr. In order that such phenomenon occurs within 100 yr, the KL time-scale should be less than 100 yr. So we investigate how many years are required for the deviation of Δ P from Δ P 1 to become large enough. For 19 200 models with different m 3 and a out , we calculate Δ P and judge the deviation simply by the relative difference between Δ P and Δ P 1 , i.e. (Δ P − Δ P 1 )/Δ P 1 > 1. We define T d as the time when this criterion is satisfied. Fig. 3 shows the colour contour map of T d . The black region in the bottom-right corner of Fig. 3 corresponds to T d > 100 yr. The white dotted line shows the critical curve used by the Newtonian formula T KL = 100 yr given by equation ( 3 ). It is also found that the dependence of T d on a out and m 3 is consistent with the Newtonian formula T KL . This is because the semimajor axis in the present model is large and 1PN effect is quite small. The top-left white region in a out 15 au and m 3 > 30M ⊙ shows that the system becomes unstable if the initial parameters are in the region. The yellow solid line is the empirical criterion for the instability given by Blaes et al. ( 2002 ): From Fig. 3 , we find that if this binary has the third companion with m 3 = 10–100M ⊙ and , we may be able to see the bending of the cumulative-shift curve within 100 yr through the observation of radio pulse. Figure 3. The parameter range of third body with which the bending of cumulative-shift curve occurs within 100 yr. The horizontal and
Paper: 2018MNRAS.481..835O (acknowledgments)
ACKNOWLEDGEMENTS The authors are grateful to the referee, Dylan Nelson, for detailed and constructive comments that enhanced this work. The authors would like to thank Joe Burchett, Ryan Horton, Ali Rahmati, John Stocke, Todd Thompson, Colin Norman, Ryan O’Leary, and Jason Prochaska for interesting conversations that contributed to this manuscript. We acknowledge discussions with Sean Johnson and Zachary Hafen about the Turnbull statistical method. Support for Oppenheimer was provided through the NASA ATP grant NNX16AB31G. This work was supported by the Netherlands Organisation for Scientific Research (NWO) through VICI grant 639.043.409. This work used the DiRAC Data Centric system at Durham University, operated by the Institute for Computational Cosmology on behalf of the STFC DiRAC HPC Facility ( www.dirac.ac.uk ). This equipment was funded by BIS National E-infrastructure capital grant ST/K00042X/1, STFC capital grants ST/H008519/1 and ST/K00087X/1, STFC DiRAC Operations grant ST/K003267/1, and Durham University. DiRAC is part of the National E-Infrastructure. RAC is a Royal Society University Research Fellow. AJR is supported by the Lindheimer Fellowship at Northwestern University.
Paper: 2021MNRAS.504.1657K (fulltext)
of destroyed GCs to field populations in the bulges of MW-like galaxies, comparing the predictions with the measurements by Schiavon et al. ( 2017 ). They show that, for most of the MW-like galaxies in their simulated volume, the prediction for f Nr of the metal-poor stellar population is lower than the observations by factors of ∼2–30 (bottom panel of their fig. 4). However, for a few simulated galaxies, the predicted f Nr are in good agreement with the observations. Like the MW, the disc populations of those galaxies are characterized by a bimodal distribution in the α–Fe plane, which is a distinctive feature of the MW disc populations (e.g. Hayden et al. 2015 ; Mackereth et al. 2017 ). Mackereth et al. ( 2018 ) showed that this feature is associated with an atypical accretion history, characterized by intense merging in early times and relative calm since z ∼ 1–1.5. It is noteworthy, however, that Hughes et al. ( 2020 ) predictions for these few MW-like galaxies differ from our measurements with regard to the dependence of f Nr on position in the α–Fe plane. The high frequency of ex-GC stars in the field of simulated galaxies is predominantly due to high-α in situ GC formation and destruction, whereas our data show that the high f Nr in the MW bulge is due to the contribution by the dissolution of low-α accreted GCs. This discrepancy would be alleviated if some of the stars in the accreted region in Fig. 2 were in fact formed in situ (see fig. 2 of Hughes et al. 2020 ), but it is not clear that accounting for such a contamination would completely eliminate the disagreement. 4 SUMMARY The results presented in this paper make use of elemental abundances from APOGEE DR16 along with data from Gaia DR2 to study the chemical and kinematic properties of 146 N-rich stars located within the inner 4 kpc of the Galaxy. Our conclusions can be summarized as follows: We find that there are likely accreted and in situ components to the N-rich population within 4 kpc of the Galactic Centre, identified via chemistry by making a cut in [α/Fe]–[Fe/H] space towards low metallicities (as shown in Fig. 2 ; e.g. Hayes et al. 2018 ; Mackereth et al. 2019 ; Das et al. 2020 ). By making this cut and removing stars without proper motions in Gaia , we select 428 and 13 982 bulge stars that lie in the accreted and in situ positions, respectively, with 25 N-rich stars being located in the accreted and 58 located in the in situ locus. We show that our sample of N-rich stars occupies the same locus as so-called SG GC stars, supporting the idea that they are the by-products of GC destruction/evaporation. We find that there is a significant difference in the kinematic properties of chemically defined accreted and in situ populations. This shows that our chemistry-based distinction of these populations maps into differences in kinematic space. We also find that the accreted bulge field population includes stars that share orbital properties with the GE/S system, although no