adsabs/WIESP2022-NER huggingface dataset:
Paper: 2020MNRAS.494.4382S (fulltext)
for σ u and bottom panels for σ l ; left panels for substituting ν + △ν and right panels for substituting ν − △ν into the MHD equations. In order to make a more reasonable estimate of the NS mass, we consider the error of QPOs. After substituting ν upper + △ν or ν upper − △ν into the MHD equations, the mass of the NS in 4U 0614+09 can be estimated as M 2.60 M ⊙ or M 2.50 M ⊙ from the minimum σ l . Similarly, after substituting ν lower + △ν or ν lower − △ν into the MHD equations, the mass of the NS in 4U 0614+09 can be estimated as M 2.50 M ⊙ or M 2.60 M ⊙ from the minimum σ u . The results are shown as panels (a), (b), (c) and (d) of Fig. 8 . Similar to the above process, the mass of the NS in 4U 1636–53 can be estimated as 2.05 M ⊙ M 2.31M ⊙ , 2.00 M ⊙ M 2.31M ⊙ with σ u from the panels of Fig. 9 . In summary, the mass of the NS in 4U 0614+09 is estimated as M 2.60 M ⊙ and that in 4U 1636–53 as 2.00 M ⊙ M 2.31M ⊙ . We do not recalculate the result of 4U 1608–52 because there are very limited results with errors reported. Figure 9. The minimum like-standard deviations σ l and σ u corresponding to the different masses of the NS in 4U 1636–53 when the errors of the observed kHz QPOs are considered. Top panels for σ u and bottom panels for σ l ; left panels for substituting ν + △ν and right panels for substituting ν − △ν into the MHD equations. Figure 10. Mass–radius relationships for various equations of state given in Özel Freire ( 2016 ). The dotted line, twin solid lines and twin dash–dot lines correspond to the masses M = 2.60 M ⊙ , , , respectively, which are obtained from the estimation by the twin kHz QPOs in 4U 0614+09, 4U 1636–53, 4U 1608–52. 5 DISCUSSION As the dividing line between NSs and black holes, EOSs affect the number of observable NSs and their properties, because different EOSs might lead to different mechanisms of supernova explosions (Özel Freire 2016 ). As seen from Fig. 10 , different EOSs have different maximum masses of NSs. Thus a high-mass NS can be used to rule out some EOSs. For 4U 0614+09, M 2.60 M ⊙ , and no constraint can be put on any EOS. For 4U 1636–53, 2.00 M ⊙ M 2.31 M ⊙ rules out EOSs of strange quark matter, which is consistent with the conclusion obtained from the lower limit on the NS radius (11.0 km) that was inferred from observed mHz QPOs modelled as marginally stable burning on the NS surface (Stiele, Yu Kong 2016 ). Güver et al. ( 2010 ) estimated the mass and radius of the NS in 4U 1608–52 when the most probable distance () was measured by the interstellar extinction towards the source. They obtained a mass of M = 1.74 ± 0.14 M ⊙ and a radius of . Steiner, Lattimer Brown ( 2010 ) obtained a larger radius of the NS (about ) when the photospheric radius is much larger than the radius of the NS. Poutanen et al. ( 2014 ) studied 21 photospheric radius expansion bursts from 4U 1608–52 and obtained a NS radius R larger than 12 km for 1.0 M ⊙ M 2.4 M ⊙ , and R restricted to be between 13 and 16 km for
Paper: 2017A&A...601A.113S (acknowledgments)
We thank the anonymous referee for useful comments and suggestions that helped to improve the manuscript. P.C. acknowledges financial support of the European Research Council (ERC; project PALs 320620). M.J. acknowledges the support of the Academy of Finland Grant No. 285769.
Paper: 2017ApJ...837...82H (fulltext)
of sight lines that cross through the z -axis of the domain behind the cloud’s head. The angles between each of these sight lines and the sight line through the cloud’s head were determined from the assumption that the observer was on Earth, 12.4 kpc ( L08 ) from the SC (see Figures 1 , 2 , and 4 ). Figure 5 (a) shows as a function of this angle for Model A at various epochs. For angles ≲30°, reaches a reasonably steady state by t = 300 Myr (i.e., the variation of from epoch to epoch is relatively small compared to earlier epochs). However, there is still some time variation of the profiles, due to the fact that the flow in the model HVC’s tail is not steady, and for the last few tens of megayears of the simulation, is drifting toward larger values for angles ≲15°. For Models B and C (Figures 5 (b) and (c), respectively), we find that also reaches a reasonably steady state by around t = 300 Myr, although again there is some time variation of the profiles. In addition, the Model C profiles exhibit gaps at some angles, meaning that there is no material above the high-velocity cutoff along the sight line at those particular angles. These gaps in the high-velocity material distribution are due to the higher-density ambient medium having a stronger braking effect on the cloud material than in the other models (this braking effect was noted in footnote 1 for a model with an even higher ambient density, for which the effect was more severe than in Model C). Figure 5. Fraction of the high-velocity material originating in the halo, , as a function of angular distance behind the cloud head, extracted from various epochs of Models A–C (top to bottom; see keys). The vertical dashed lines indicate the location of F16 's sight lines relative to the SC head. Figure 6 shows averaged over multiple epochs. Within the epoch-to-epoch variations, these averaged profiles are similar for Models A and B at the cloud head and at angles ≳20° behind the cloud head. The difference between the models at intermediate angles is because the Model A cloud fragments more than the Model B cloud (as noted above), leading to more rapid mixing in the first few kiloparsecs behind the cloud head. However, other than in the first few kiloparsecs behind the cloud head, the mixing results are not very sensitive to the assumed cooling rate. In contrast, the averaged profiles from Models A and C are similar at all angles, though this is in part due to the large time variation of the Model C profiles. These results suggest that the extent of mixing is not very sensitive to the assumed ambient density. Figure 6. sampled every 1 Myr and averaged over t = 300–400 Myr from Models A, B, and C (red, blue, and green, respectively). The colored bands indicate the standard deviation. The vertical dashed lines indicate the location of F16 's sight lines relative to the SC head. Figure 7 compares the averaged profiles calculated from versions of our reference model with different maximum resolutions.
Paper: 2020MNRAS.496.1445C (fulltext)
II stars (as modelled in Delphi ). Indeed, a larger contribution from Pop III stars would lead to a faster coupling between the spin and CMB temperatures, resulting in a higher redshift at which the signal would tend to 0. Finally, in all panels, we also show by the grey shaded regions the models that are consistent with the EDGES data. As expected, while the range of histories naturally shows a dispersion, their redshift-evolution remains unaffected. The Pop III SFRD and the radio emissivity show the least dispersion (since they effectively govern the brightness temperature from Pop III stars) while, given their negligible contribution to the X-ray emissivity, the value of shows the largest dispersion. It is worth pointing out here that the depth of the EDGES signal is highly controversial (see, e.g. Hills et al. 2018 ; Sims Pober 2020 ). In case the amplitude of the signal turns out to be much smaller than what is assumed in this work (say |δ T b | ≲ 100 mK), then the Pop II stars with conventional choice of efficiency parameters are sufficient to explain the observations. Hence our conclusions on the Pop III stars hold only in the case where the amplitude and redshift extent of the 21-cm signal are similar to those reported by the EDGES experiment. Finally, we compare our predicted redshift evolution of the Pop III SFRD with some theoretical models in Fig. 3 . Most models predict a Pop III SFRD that either rises between z ∼ 22 and 13 (Jaacks et al. 2018 ; Sarmento et al. 2019 ) or shows a relatively constant behaviour for z ∼ 20 − 13 (Johnson et al. 2013 ). The only exceptions are the predictions by Valiante et al. ( 2016 ), Valiante et al. ( 2018 ), whose amplitude and redshift dependence are very similar to our predicted shape. Interestingly, the Valiante et al. ( 2016 ), Valiante et al. ( 2018 ) models aim to describe the formation histories of z ∼ 6–7 quasars and their host galaxies. They therefore describe rapidly evolving and biased regions of the Universe. In these models, Pop III stars (and their black hole remnants) start to form at z > 20 in dark matter mini-haloes with mass [10 6 –10 7 ]M ⊙ . However, their formation efficiency is small, as Lyman-Werner radiation limits gas cooling due to H 2 photo-dissociation. Because of this, the bulk of Pop III stars form at redshift 15 z 20 in metal poor ( Z Z cr = 10 −3.8 Z ⊙ ) Ly α-cooling haloes, with mass [10 7.5 − 10 8.5 ] M ⊙ that are able to self-shield and cool the gas more efficiently (see the top panel of Fig. 1 in Valiante et al. 2018 ). At z ≤ 15 Pop III star formation is suppressed by metal enrichment, as Z Z cr = 10 −3.8 Z ⊙ star forming regions become increasingly rare. Hence, the shape of the EDGES signal suggests that the rapid coupling between the spin temperature and the gas kinetic temperature is driven by Ly α photons produced by Pop III stars that form in relatively massive and rare haloes, whose abundance grows fast, due to the exponential growth of the massive-end
Paper: 2017MNRAS.469.1360C (acknowledgments)
Acknowledgments This paper is based on observations collected at the La Silla Observatory (ESO, Chile) with the FEROS spectrograph at the 2.2-m telescope under programme 088.C-0892(A). This paper includes data collected by the K2 mission. Funding for the K2 mission is provided by the NASA Science Mission directorate. The authors acknowledge the support of the UK Science and Technology Facilities Council (STFC). Funding for the Stellar Astrophysics Centre is provided by The Danish National Research Foundation (Grant DNRF106). DV received funding from the European Research Council under the European Unions Seventh Framework Programme (FP/20072013)/ERC Grant Agreement n. 320964 (WDTracer). TM acknowledges financial support from Belspo for contract PRODEX GAIA-DPAC. DH acknowledges support by the Australian Research Councils Discovery Projects funding scheme (project number DE140101364) and support by the National Aeronautics and Space Administration under Grant NNX14AB92G issued through the Kepler Participating Scientist Programme. MNL acknowledges the support of The Danish Council for Independent Research Natural Science (Grant DFF-4181-0415). TSR acknowledges the support from CNPq-Brazil and PRIN INAF 2014 CRA 1.05.01.94.05. This research has made use of the SIMBAD data base, operated at CDS, Strasbourg, France.
Paper: 2018ApJ...868...60G (fulltext)
different ZDI observations of the same star. We compare the complexity derived from the observations to the magnetic complexity function in CG18 , , in Figure 1 . Here, Ro is the commonly used Rossby number for stellar magnetic dynamo activity representing the ratio Ro = τ c / P rot of rotation period, P rot , and convective turnover time, τ c . This complexity function was derived by matching the evolution of rotation periods of stars in young open clusters using the relationship between magnetic complexity and angular momentum loss derived by CG16 . The function shown in Figure 1 that provides a good match to the data corresponds to a convective turnover time τ c = 100 and value a = 0.005. The former is appropriate for the spectral types of the secondary stars considered here (Morin et al. 2010 ; Wright et al. 2011 ). The value of a is four times smaller than the one used by CG18 , though we note that this quantity only affects the most rapidly rotating stars with the smallest Rossby numbers and was not well constrained in that study that only dealt with stars with masses greater than 0.3 M ⊙ . In our CV orbital evolution model (Section 3 ), we adopt an intermediate value a = 0.01. The shape of the complexity function, n , which dictates the change in angular momentum loss with changing rotation period, remains the same. It should also be kept in mind that, while n av was shown to work reasonably well for parameterizing real magnetograms (Garraffo et al. 2016 ), there is some arbitrarity in its definition and it should strictly be interpreted only as a relative measure of complexity. The spatial resolution of ZDI-reconstructed magnetic fields is also limited by the rotation-induced velocity shift of surface magnetic features and the deconstructed complexity is inevitably going to be a lower limit to the true complexity. 3. Magnetic Disruption of Angular Momentum Loss 3.1. Angular Momentum Loss Here, we simulate the orbital evolution of CVs, including, for the first time, the magnetic complexity modulation and its effect on angular momentum loss. We use the same spin-down model as presented by CG18 . As systems evolve and lose angular momentum, the orbits shrink and periods decrease. Efficient spin–orbit coupling means that the secondary stars spin up, maintaining synchronicity between their rotation and orbital periods. Figure 1 indicates that their magnetic complexity will also increase as they spin up. The data show a somewhat steeper increase in complexity for the stars with the shortest rotation periods, although data are lacking for periods less than 10 hr. We proceed by adopting the CG18 spin-down model and complexity function, taking Figure 1 as offering some empirical support (with the function reproducing well the data on average), but recognizing that it still represents somewhat of an extrapolation to reach the CV period gap. The spin-down model can then be described by the following equations (see CG18 for further details):
Paper: 2017MNRAS.466.4530C (fulltext)
in the two transitions. There is no known maser association with this core but the two are both associated with H 2 O (clusters 1 and 2, respectively), and CH 3 OH masers (Slysh etal. 1999 ). Based on our detections of two bipolar outflows in Sh2-76E, we can confirm the presence of at least two YSOs. This does not rule out the possibility of the presence of a much earlier stage YSO that is not yet driving maser emissions. However, we have explore the infrared point source catalogue to verify the nature of the possible driving sources of the outflows. We have selected infrared source with 1arcsec from the cluster positions and used them to construct the spectral energy distribution (SED) of the driving sources using Robitaille etal. ( 2006 ) SED fit tool while adopting our parallax distance. We found Two Micron All Sky Survey (2MASS) and Wide-field Infrared Space Exploration ( WISE ) points sources associated with Sh2-76EMM1 but only 2MASS point source for Sh2-76EMM2. Fig. 6 shows the SED of Sh2-76EMM1 (left-hand panel) and Sh2-76EMM2 (right-hand panel). The SED of Sh2-76EMM1 corresponds to that of a class I YSO. In the case of Sh2-76EMM2, only three data points from 2MASS was used and the error in the fitting is large. We cannot conclude on the class of the YSO/driving source Sh2-76EMM2. However, the 2MASS K -band and H -band magnitude were very similar, thus we speculate that the YSO could be a class II object. Figure 6. P(top) SED of Sh2-76EMM1 and (bottom) SED of Sh2-76EMM2. 4.2 A Plausible Probe for Binary/Multiplicity Along Same Line of Sight Line-of-sight effect on the observed properties of an object can be influenced strongly by the interception effect of another object. This could be in the form of strong absorption of the molecular line emission from the background source. With a reasonable angular resolution, the nature of molecular line emission could reveal the presence of more than one source within a region. The two driving sources of Sh2-76Eare located along the same line of sight with V LSR difference of 5kms 1 . Zinchenko etal. ( 1994 ) reported an unusual HCN emission spectrum with a missing main component, which they argued may have been caused by a foreground, 5kms 1 blue-shifted cloud. Zinchenko etal. ( 1997 ) also reported two core components based on their NH 3 (1,1) spectra, which showed additional component after a single Gaussian fitted component was subtracted. In the case of the HCN emission, it is imperative that only a very dense core could have completely absorbed the main component of the line emission. Self-absorption could have affected only the emissions around the systemic velocity of the central source, leaving a fraction of the emission to be detected. The NH 3 line observational results of Zinchenko etal. ( 1997 ) support the existence of a second component. We propose the Sh2-76EMM2 is responsible for the absorption of the HCN main component. MM2 is only 0.5arcsec from MM1 and will provided the
Paper: 2017A&A...605A.128K (fulltext)
indicated by an arrow. We note, however, a group of six stars located within 2.6 core-radii of the center (red circles in Fig. 2 ) that show very similar kinematic properties, clustering around a mean velocity of km s. Likewise, as we will show in Sect. 4, their the metallicities are in excellent agreement with each other. While the presence of a secondary plume of stars at km s could in principle be identified with an underlying cluster population as well, these stars have markedly different metallicities (Sect. 4) and we rather deem these foreground objects. The same holds for those stars outside of the cluster boundary of , Thus we identify the innermost six stars at the higher velocity peak with members of FSR 1716. Their individual properties are listed in Table 2 . Table 2 Derived properties of the FSR 1716 member candidates. Thereby, we establish the mean systemic velocity of FSR 1716 as km s with a velocity dispersion of km s, a moderately low value typical of a low-to-intermediate mass Galactic star cluster (e.g., Pryor Meylan 1993 ; Koch et al. 2012 , 2017b ). Using this information and the cluster morphology in terms of its characteristic King (1966) radius ( Froebrich et al. 2008 ; Bonatto Bica 2008 ; Minniti et al. 2017 ), we can estimate the dynamical mass of FSR 1716 as ( Spitzer 1987 ), which is a factor of 70 larger than the low-mass estimate of Bonatto Bica (2008) , albeit based on their premise of a significantly lower distance. Upon a typical stellar mass-to-light ratio of 1.5, this would place FSR 1716 at an approximate absolute magnitude of mag, without doubt in the regimen of a GC. Ideally, the kinematic information should be coupled with accurate proper motions, allowing for an orbital computation that could hint at this cluster’s origin. Unfortunately, none of the member stars had a cross-match in the UCAC5 catalog ( Zacharias et al. 2017 ), and only three stars could be identified in the Hot Stuff for One Year (HSOY) proper motion catalog ( Altmann et al. 2017 ), which used the most recent positions from Gaia -DR1 ( Prusti et al. 2016 ). However, the respective proper motions are highly inconclusive, leading to a broad variety of possible orbits, from thick disk to outer halo; therefore, we do not pursue the orbital analyses any further. 4. Calcium triplet metallicities The three near-infrared CaT lines are a powerful indicator of stellar metallicities, [Fe/H], for RGB stars near and above the horizontal branch ( Armandroff Zinn 1988 ). In the following, we adopted the calibrations of line strength with metallicity from Vásquez et al. (2015) , which are applicable into the metal-rich regime up to 0.7 dex, and which employ the infrared magnitude above the red clump (). The latter is highly advantageous in population studies towards low latitudes, compared to the traditional use of the visual magnitude above the horizontal branch (), which is more sensitive to reddening. In practice, these calibrations are
Paper: 2019MNRAS.486L..85C (acknowledgments)
ACKNOWLEDGEMENTS This work was supported by a grant from the Research Grants Council of the Hong Kong Special Administrative Region, China (Project No. EdUHK 28300518).
Paper: 2020MNRAS.499.2534K (fulltext)
But, under a simplistic steady state assumption, our results do not favour the surface density thresholds required for star formation commonly adopted in Galactic studies that assume a certain magnetic field strength (Heiderman et al. 2010 ). However, low-density gas in the Magellanic Bridge has been evolving into denser molecular gas, from which stars form, and the lower densities of the molecular gas in this region may signal the exhausting of the latter. The density of YSOs and PMS stars in the subregions ranges in a continuum from moderately clustered to relative isolation, reminiscent of star formation in the Milky Way (e.g. Bressert et al. 2010 ). However, low number statistics precludes us from defining a surface density for each subregion. Our results suggest that star formation observed in MB -C is underway in localized regions depending on the density inhomogeneities within an initial cloud, and is intimately linked with the presence of dense molecular gas. 5.2 Molecular gas structure in the Magellanic Bridge In Section 2 , we have investigated the properties of the molecular clumps in MB -C. The clumps were gravitationally stable, and very similar to clumps related to star formation in the Milky Way. Their densities, radii, masses, and linewidths do not differentiate them from clumps previously identified in the intermediate to low-density regimes of star formation in the Milky Way (Kennicutt Evans 2012 ). The clumps are related to star formation, as seen by their proximity to YSOs and PMS stars. The properties of YSOs and PMS stars show no remarkable differences compared to the Galaxy on account of their metallicities, except lower extinctions towards them. Overall, this suggests that even though the CO has lower filling factors in the Magellanic Bridge than found in solar metallicity molecular clouds, the actual star formation process proceeds similarly once cold molecular gas is formed, and differences may arise due to the conversion of atomic to molecular gas. The lower dust-to-gas ratio results in lower extinction towards the line of sight of protostars, which affects similarly the X CO factor, but the formation of molecular gas is the unifying factor between star formation. The effect of the lower self-shielding by CO on its durability is unclear, since the incident radiation field in MB -C is not particularly strong only containing early B-type stars. Recent observations from the Herschel space telescope at far-infrared wavelengths have suggested a common unifying factor between star formation and the formation of filamentary structure in the cold interstellar medium (André 2017 ), ubiquitous in star-forming regions. Unfortunately, the FWHM of Herschel imaging (∼5–10 pc) at the distance to MB -C is more than an order of magnitude larger than the average width of filaments detected in the Milky Way (∼0.1 pc) therefore it is currently not possible to identify filaments in the Magellanic Bridge. Instead, we simply overplot YSOs
Paper: 2016MNRAS.462.3243Z (acknowledgments)
We thank the anonymous referee for helpful comments and suggestions. This work is partly supported by the National Natural Science Foundation of China (Grant No. U1431126; 11263002; 11311140248; 11203026) and Guizhou Natural Science Foundations (201519; 20134021; 20117006). HYC was supported by BK21 Plus of the National Research Foundation of Korea and a National Research Foundation of Korea Grant funded by the Korean government (NRF-2015M1A3A3A02009155).