adsabs/WIESP2022-NER huggingface dataset:
Paper: 2015MNRAS.452.4252G (acknowledgments)
We would like to thank Dr Marco Milla, the director of Radio Observatorio de Jicamarca, for his help of initiating this experiment and collecting data. We thank JRO staff for performing the observations. We also thank Karim Kuyeng and Ramiro Yanque for clarifying many questions regarding these observations. JRO is a facility of the Instituto Geofisico del Peru operated with support from the National Science Foundation (NSF) grant AGS-0905448 through Cornell University. This effort was supported under NSF grants ATM 07-21613 and AGS 12-02019 to Pennsylvania State University.
Paper: 2017ApJ...837..137M (acknowledgments)
We are grateful to the anonymous referee and to Junichi Baba, Tomoki Morokuma, Shinya Komugi, Robert C. Kennicutt, Daisuke Iono, Tetsuo Hasegawa, Tsutomu Takeuchi, Kouichiro Nakanishi, Fumi Egusa, and Nario Kuno for enlightening comments and suggestions. We thank Eric Feigelson for his constructive comments on the statistical aspects. K.M. thanks Ross Burns for the English language editing.
Paper: 2017ApJ...843..143S (acknowledgments)
The Liverpool Telescope is operated on the island of La Palma by Liverpool John Moores University in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias with financial support from the UK Science and Technology Facilities Council (STFC). C.G.M. acknowledges support from the Royal Society, the Wolfson Foundation, and STFC. D.K. acknowledges the financial support from the Slovenian Research Agency (research core funding No. P1-0188).
Paper: 2016ApJ...832..131P (acknowledgments)
The authors would like to thank Heather Knutson for helpful discussions throughout the preparation of this manuscript. 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. 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. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation. This work was partially supported by funding from the NSF Astronomy Astrophysics and NASA Exoplanets Research Programs (grants AST-1109857 and NNX16AI14G, G.A. Blake P.I.), and the Center for Exoplanets and Habitable Worlds, which is supported by the Pennsylvania State Unviersity, the Eberly College of Science, and the Pennsylvania Space Grant Consortium. Basic research in infrared astrophysics at the Naval Research Laboratory is supported by 6.1 base funding. Finally, we thank an anonymous reviewer for helpful insights which improved the content of this paper.
Paper: 2020A&A...643A.134B (fulltext)
represents an effect of the detailed distribution of the magnetic field over the visible hemisphere of the MWD as viewed during our second observation. It is very interesting to consider our data for WD 1703−267 in the light of our earlier effort (see Sect. 3 ) to calibrate the order-of-magnitude relationship of Eq. ( 1 ). For this star, we have accurate values of both ⟨| B |⟩ and ⟨ B z ⟩, and judging from the simple V ∕ I profiles through the Balmer lines, a relatively simple magnetic field structure. With ⟨ B z ⟩≈ 2.75 MG in our second observation, we predict from our calibration of the coefficient of Eq. ( 3 ) that there should be continuum circular polarisation of V ∕ I ≈ 0.2%. Rebinning our data over 16 pixels to improve the S/N, we find that a weak continuum circular polarisation is indeed present at the level of V ∕ I ≈ 0.25% (see Fig. 4 ). Furthermore, the relationship between the sign of the continuum polarisation and that of V ∕ I in the blue σ components of the Balmer lines is the same as in G99-47. This result reinforces confidence in the use of broad-band circular polarisation as a tool for measuring ⟨ B z ⟩, and in the general accuracy of our calibration. 5.3 DA star WD 0810−353 = UPM J0812−3529 This star was identified as a WD within the 20 pc volume by the USNO parallax programme ( Finch et al. 2018 ), and is included in the list of WDs within the 20 pc volume established by Hollands et al. (2018) on the basis of the Data Release 2 dataset ( Gaia Collaboration 2018 ). For the reasons discussed in connection with WD 1703−267, we provide an (unofficial) Villanova WD number for the star, which we use throughout. Physical parameters based on Gaia astrometry and photometry ( Hollands et al. 2018 ; Gentile Fusillo et al. 2019 ) are provided in Table 1 . The spectral class is unknown, asno previous spectra of the star appear to have been taken. As the effective temperature is high enough that we would expect H β to be observable if the atmosphere of this object is H-rich, while possible metal lines around 4000 Å would be of great interest, we obtained circular spectropolarimetry of this star with the 600B and 1200B grisms. WD 0810−353 is bright enough that we were able to obtain quite high S/N (450–500 per Å). Fig. 4 WD 1703−267 (observations of September 30, 2019): circular polarisation (red solid line) and null profile (blue solid line, offset by − 3%), superposed to the V ∕ I error bars (light blue). Spectra are rebinned at about 11 Å. The three Stokes V ∕ I spectra are in quite good agreement in that they all show a region of weak positive polarisation around 4720 Å, and two regions of stronger negative polarisation around 4830 and 4960 Å. The I spectra obtained with grism 1200B show a series of weak depressions (with an amplitude of the order of 1%) that are hard to disentangle from instrumental features, even after applying flat-fielding correction. After an approximate normalisation to the continuum (clearly difficult to
Paper: 2019MNRAS.484.4347V (fulltext)
Note the systemic velocity is relative to the template file used. We do note the systemic velocity we find, Γ = 321.18 ± 0.85 km s −1 , is higher than the 295.8 ± 2.0 km s −1 previously reported (Corbet et al. 2016 ). We undertook additional analysis to confirm that the wavelength calibration performed by the HRS pipeline was correct and that the comparison to the radial velocity standards was accurate to within the expected performance. We found no evidence of any discrepancy in the calibration, nor any long-term systematic shift in the wavelength calibration over the period of observations. We believe that there are two possible reasons for this difference; there may be a possible systematic offset arising from the fit of the high-resolution template fit to the medium resolution MILES libraries (however there is no significant offset to the radial velocity standards) or there may have been a systematic offset in the zero velocity in the field O-type star used in the previous analysis. However, this difference does not change the main results of determining the orbital parameters. If this difference in the systemic velocity is removed the radial velocities are in agreement with the previous results. The final radial velocity curve and the best-fitting model are shown in Fig. 2 , and the data are given in Table 3 . We find the binary is slightly eccentric, e = 0.40 ± 0.07, and place the time of periastron at HJD 2457412.13 ± 0.29. Figure 2. Radial velocity of the O5 iii star in the binary system, as determined from cross-correlation. The solid line shows the best fit to this data, with a fixed 10.301 d orbital period. For clarity the plot is reported over two orbital phases with phase φ = 0 at MJD = 57 410.25, which corresponds to the given phase in Corbet et al. ( 2016 ). The error bars show the statistical errors reported by the rvsao package. Table 3. Radial velocities obtained from cross-correlation. 3 DISCUSSION 3.1 Mass of the compact object The binary parameter solution gives a mass function of f = 0.0010 ± 0.0004 M and the constraints on the mass of the compact object are shown in Fig. 3 . The mass of an O5 iii (f) star is ∼40 M (Martins et al. 2005 ), but due to the binary evolution could be different. A mass range of 25–42 M was considered in Seward et al. ( 2012 ). For this mass range both a neutron star or black hole mass is compatible with the mass function, though a neutron star is favoured. For a 40 M star, the mass of the compact object will be >5 M for inclinations i ≤ 15 ± 2°, and while the inclination must be i ≤ 11 ± 1° for a 25 M optical companion. Assuming the compact object is a pulsar, with a mass of 1.4 M , the inclination will lie between i = 39 ± 6° and i = 59 ± 11° for a 25 M and 40 M star, respectively. Figure 3. The constraint on the mass of the compact object in LMC P3. The shaded area marks the range of assumed masses of the optical companion. 3.2 Binary orientation The orientation of the binary system is
Paper: 2020MNRAS.491.4960L (acknowledgments)
ACKNOWLEDGEMENTS This work is supported in part by the National Natural Science Foundation of China under Grant No. 11675032 (China).
Paper: 2020MNRAS.496L..43C (fulltext)
luminosity wind that emanates from a stellar surface or a galaxy, the interaction between the wind supplying the energy and the ambient gas forms a ‘bubble’. This bubble consists of a forward shock that advances into the ambient gas, a reverse shock that travels back through the wind, and a contact discontinuity that separates the two shocked fluids (e.g. Castor, McCray Weaver 1975 ; Weaver et al. 1977 ; Fielding, Quataert Martizzi 2018 ); a self-similar solution describes the fluid properties between the forward shock and the contact discontinuity. When the shock generated from a supernova explosion or an expanding wind is young, the large temperatures and densities at the shock front imply that the effective adiabatic index γ of the multispecies fluid satisfies 4/3 ≲ γ ≲ 5/3. As the shock further expands and cools, however, the high-density regions near the shock front start to radiate efficiently, which reduces γ below the value of an ideal monatomic gas. On the other hand, the low-density material in the interior of the explosion remains hot, and the expanding blast wave is characterized by a thin, nearly isothermal shell that is in contact with a hot, pressurized interior (e.g. McKee Ostriker 1977 ; Shull McKee 1979 ; Blondin et al. 1998 ). Once the shock becomes radiative, it is no longer strictly correct to treat the fluid as adiabatic. However, it is reasonable to assume that the gas exhibits an increased degree of compressibility prior to radiating a significant amount of its total energy, during which time it is still approximately correct to use the self-similar scalings (derived under strict adiabaticity) for the relationship between the shock position and velocity, but with an appropriately reduced adiabatic index. At the very least, the self-similar solutions with a nearly isothermal equation of state describe the qualitative behaviour we expect at this stage – a nearly evacuated, extremely hot interior in contact with a thin shell that contains the majority of the mass (e.g. Fig. 1 ). Understanding the structure of these shocks therefore yields insight into the more general behaviour of radiative, thin shells formed from explosions and stellar and galactic winds. Figure 1. A comparison of the exact solution for the Sedov–Taylor blast wave (solid lines) and the analytical, leading-order solution in the quantity δ = (γ − 1)/(γ + 1) given by equation ( 13 ) (dashed lines) for the adiabatic indices γ and radial power-law indices n shown in the legend; recall that the ambient density follows the power-law decline ρ ∝ R − n , where R is the shock position. Solutions are plotted as functions of the standard Sedov–Taylor variable ξ = r / R ( t ), which is just the spherical radius at a given time normalized by the shock radius. The left, middle, and right columns show the radial velocity normalized by the shock velocity, the density normalized by the ambient density (and reduced by the factor δ), and the pressure normalized by the
Paper: 2016A&A...589A..74D (acknowledgments)
This work is based on observations carried out with the IGN 40 m radio telescope. We thank all the staff in Yebes Observatory for the generous help received. We acknowledge partial funding from MINECO grants AYA2012-32032, FIS2012-32096, and FIS2012-38160. This work has made use of the American Association of Variable Star Observers (AAVSO) database and of the SIMBAD Astronomical database (operated at CDS, Strasbourg, France).
Paper: 2021ApJ...913...92R (fulltext)
2009 , 2010 ; Lyra Klahr 2011 ; Lyra 2013 ). Yet it remains 2–3 orders of magnitude lower than the radial and azimuthal velocity components Figure 3 . Figure 1. Vertical vorticity of unstratified gas in a three-dimensional box at different times. The convective overstability amplifies seed noise into growing vortices, that intensify and merge viscously, saturating at radial length similar to the pressure scale height. Figure 2. The vertical component of the vorticity () in the midplane (left panels) and a vertical cut through the disk at the azimuthal location of maximum negative vorticity (indicated by the dashed line). Also shown is the vertical velocity component ( u z ). The upper plots are shown at 100 local orbits (top panels). At this time, there is no structure in the vertical direction: the flow is essentially two-dimensional. After 275 local orbits (bottom panels) there is now turbulent motion inside the vortex, which is also visible in the vertical velocity. This is an indication of the growth of elliptical instability in the vortex cores. Figure 3. Gas velocity components averaged over the entire box. As the elliptical instability sets in, the vertical velocity component (green line) increases steeply, while staying well below the azimuthal and radial components. 4.1. Pebble Evolution We first discuss the run without baroclinic driving. As expected, the pebbles settle to the midplane, soon triggering the streaming instability. Although we assumed an overall initial dust-to-gas ratio of , the initial dust-to-gas ratio in the midplane is already and increases even further as the sedimentation progresses. The maximum dust-to-gas ratio of the simulation with pebbles and without baroclinic driving can be seen in Figure 4 (red line). The other two lines in the plot show the maximum dust-to-gas ratios for the simulation with , (green line), and pebbles (black line). Figure 4. Maximum dust-to-gas ratio for simulations with (black line) and (green and red line) pebbles. indicates that no baroclinic feedback was included in the simulation. Although the high pebble concentration in the simulations with baroclinic effects disrupts the vortex in the midplane, where pebbles are located, the maximum dust-to-gas ratio is higher than without baroclinic effects. The maximum concentration for with baroclinic driving is only a factor of 2 higher than in the case without it. Inspecting the evolution of the dust density in the baroclinic simulation, we see that, as expected, the pebbles settle to the midplane and get trapped in the vortices. Yet, after only three orbits, the pebbles start to perturb the vortices in the midplane via the streaming instability. Like in the 2D simulations of Raettig et al. ( 2015 ), high pebble densities generate local perturbations of the velocity field visible as steep vorticity gradients. Figure 5 shows the vorticity at the top of the box (left panel) and in the midplane (middle panel), and the dust-to-gas ratio in the
Paper: 2015A&A...579A...6S (acknowledgments)
We thank Michele Bellazzini for a very careful reading and useful comments and suggestions that helped to improve the quality of the paper. We also thank the anonymous referee for comments that greatly improved this paper. Based on data products from observations made with ESO Telescopes at the La Silla Paranal Observatory under program ID 188.B-3002. These data products have been processed by the Cambridge Astronomy Survey Unit (CASU) at the Institute of Astronomy, University of Cambridge, and by the FLAMES/UVES reduction team at INAF/Osservatorio Astrofisico di Arcetri. These data have been obtained from the Gaia -ESO Survey Data Archive, prepared and hosted by the Wide Field Astronomy Unit, Institute for Astronomy, University of Edinburgh, which is funded by the UK Science and Technology Facilities Council. This work was partly supported by the European Union FP7 program through ERC grant number 320360 and by the Leverhulme Trust through grant RPG-2012-541. We acknowledge the support from INAF and Ministero dell’ Istruzione, dell’ Università’ e della Ricerca (MIUR) in the form of the grant “Premiale VLT 2012”. The results presented here benefit from discussions held during the Gaia -ESO workshops and conferences supported by the ESF (European Science Foundation) through the GREAT Research Network Program. I.S.R. gratefully acknowledges the support provided by the Gemini-CONICYT project 32110029. C.M. acknowledges the support from CONICYT-PCHA/Doctorado Nacional/2014-21141057. D.G. gratefully acknowledges support from the Chilean BASAL Centro de Excelencia en Astrofísica y Tecnologías Afines (CATA) gran PFB-06/2007. S.V. gratefully acknowledges the support provided by Fondecyt reg. 1130721. A.K. and N.K. acknowledge the Deutsche Forschungsgemeinschaft for funding from Emmy-Noether grant Ko 4161/1. U.H. acknowledges support from the Swedish National Space Board (SNSB). S.G.S acknowledges the support from the Fundação para a Ciência e Tecnologia, FCT (Portugal) and POPH/FSE (EC), in the form of the fellowships SFRH/BPD/47611/2008.
Paper: 2021A&A...646A.138I (fulltext)
values similar to the ratios in local SFGs with high EW(H α ). This increase results in high EW([O III ] 5007) > 1000 Å for bright z > 7 galaxies ( Roberts-Borsani et al. 2016 ). The properties of the high- z SFGs discussed above contrast with those of the z ∼ 0 main-sequence SFGs. The latter exhibit, on average, a much lower level of star formation activity and, thus, they are characterised by much lower SFRs, sSFRs, and EWs of emission lines at a fixed stellar mass. However, there is a small fraction of low- z SFGs with compact structure and other properties that are similar to those of high- z galaxies. Subsets of these compact star-forming galaxies (CSFGs) are variously called blue compact dwarf (BCD) galaxies, ‘green pea’ (GP) galaxies, and luminous compact galaxies (LCGs) in the literature. BCDs are nearby dwarf galaxies at redshift z typically ≤0.01 − 0.02, with strong bursts of star formation and low oxygen abundances of 12 + logO/H ≲ 7.9, extending down to 12 + logO/H ≲ 7.0 (e.g., Thuan Martin 1981 ; Izotov et al. 1994 , 2018c ). Typically, their stellar masses are low, M ⋆ ≲ 10 8 M ⊙ . A subset of the lowest-mass BCDs with M ⋆ ∼ 10 6 –10 7 M ⊙ and very high O 32 ≳ 10 has been named ‘blueberry’ galaxies by Yang et al. (2017) because of their compact structure and intense blue colour on the Sloan Digital Sky Survey (SDSS) composite images. Cardamone et al. (2009) considered ‘green pea’ (GP) galaxies in the redshift range of z = 0.112–0.360, with unusually strong [O III ] λ 5007 emission lines, similar to high- z UV-luminous galaxies, such as Lyman-break galaxies and LAEs, with sSFRs up to ∼10 Gyr −1 . The green colour of these galaxies on composite SDSS images is due to strong [O III ] λ 5007 line emission (with EWs up to 1000 Å), redshifted into the SDSS r -band. Izotov et al. (2011) studied a much larger sample of 803 star-forming luminous compact galaxies (LCGs) in a larger redshift range z = 0.02–0.63, with global properties similar to those of the GPs. The sSFRs in LCGs are extremely high and vary in the range ∼1–100 Gyr −1 , comparable to those derived in high- z SFGs. Similar high sSFRs were also found for low- z LyC leaking galaxies ( Izotov et al. 2016a , b , 2018a , b ; Schaerer et al. 2016 ) and some other low- z SFGs (e.g., Senchyna et al. 2017 ; Yuma et al. 2019 ). Izotov et al. (2016c) found sSFRs of up to 1000 Gyr −1 in a sample of ∼13 000 CSFGs selected from the SDSS. Compared to BCDs, both GPs and LCGs are characterised by higher stellar masses ∼10 9 M ⊙ and higher oxygen abundances ∼7.8–8.3. Izotov et al. (2015) selected ∼5200 low-redshift (0 z 1) CSFGs and studied the relations between the global characteristics of this sample, including absolute optical magnitudes, SFRs, stellar masses, and oxygen abundances. They found that for all relations, low- z and high- z SFGs are closely related, indicating a very weak dependence of metallicity on stellar mass, redshift, and star formation rate. This finding favours the
Paper: 2018ApJ...854..145M (fulltext)
the BT-SETTL model grid (Allard 2014 ), where 100,000 linear combinations of three synthetic spectra were used to find the best-fit linear combination. The final parameters of the best-fit model to the MDM spectrum were T eff = 3595K, log g = 4.9 dex, and [Fe/H] = −1.1 dex. For the STIS spectrum, we found T eff = 3500 K, log g = 4.83 dex, and [Fe/H] = −1.3 dex. We trust that the temperature and gravity measurements from the model fit to the broad optical spectra because the method is well-calibrated throughout the subdwarf sequence (Lépine et al. 2007 ). 4. Metallicity. We calculated [Fe/H] for Wolf 1130A using the IRTF/SpeX spectrum and following the techniques outlined by Mann et al. ( 2013a ). That paper presented empirical relations between strong atomic lines in near-infrared M dwarf spectra and overall metallicity. These relations were calibrated using wide binaries containing an FGK primary and an M dwarf companion under the assumption of identical metallicities between binary components. In this work we adopt the mean of the H - and K -band relations [Fe/H] = −0.70 ± 0.12 as the metallicity for Wolf 1130A. Uncertainties account for Poisson errors in the spectrum as well as the scatter in the Mann et al. ( 2013a ) calibrations. The metallicity values in the literature have a large scatter but are all significantly subsolar. As shown by Neves et al. ( 2012 ), most methods for determining metallicity diverge below −0.6 dex. Rojas-Ayala et al. ( 2012 ) derive [M/H] = −0.45 ± 0.12 and [Fe/H] =−0.64 ± 0.17 from K -band indices. Other [Fe/H] calculations include −0.80 dex (Woolf et al. 2009 ), −0.87 dex (Stauffer Hartmann 1986b ), −0.89 dex (Bonfils et al. 2005 ), and −1.02 dex (Schlaufman Laughlin 2010 ). Schmidt et al. ( 2009 ) compiled [Fe/H]= −0.62 ± 0.10 and [Ti/H] = −0.22 ± 0.09 abundances for Wolf 1130A from Woolf Wallerstein ( 2006 ), who used the MOOG spectral synthesis software (Sneden 1973 ). Additionally, Schmidt et al. ( 2009 ) derived [O/H] = −0.45 ± 0.11 from TiO lines in R ∼ 30,000 optical spectra. Overall, our methods are most similar to those of Rojas-Ayala et al. ( 2012 ) and Schmidt et al. ( 2009 ) and we include their measurements in Table 4 . Employing the current version of the MOOG synthesis software with a Python wrapper, 14 14 https://bitbucket.org/madamow/pymoogi.git , written by Monika Adamow. we measured abundances in the IGRINS infrared spectrum. For this synthesis we used the parameters already identified above ( T eff = 3600 K, log g = 5, [M/H] = −1, and microturbulence = 1.0 ) to select similar atmosphere models. Our synthetic spectra employed Kurucz ( 1979 , 2011 ) models that were broadened to the instrument resolution and v sin i = 15.0 and then modulated for different atomic abundances. Figure 9 shows the K -band Ca i triplet along with synthetic abundances of [Ca/H] = −0.2 ± 0.1 dex. The 10% depths of the Ca lines, relative to the continuum, are some of the deeper lines in the infrared spectrum of
Paper: 2018ApJ...859..160W (acknowledgments)
We thank the anonymous referee for helpful suggestions. This work is supported by the National Basic Research Program (973 Programme) of China (grant No. 2014CB845800), the National Natural Science Foundation of China (Grants 11673006, 11533003, 11303005, 11851304, 11773007, U1731239), the Innovation Team and Outstanding Scholar Program in Guangxi Colleges, the One-Hundred-Talents Program of Guangxi colleges, and the Guangxi Science Foundation (2016GXNSFFA380006, AD17129006, 2016GXNSFDA380027).