adsabs/WIESP2022-NER huggingface dataset:
Paper: 2020MNRAS.494.5239D (fulltext)
collisions with para -H 2 ( j 2 = 0). We also observe that the relative magnitudes of the rate coefficients for the transitions to the 1 01 level are different in Figs 2(a) and (b) . Figure 2. Rate coefficients as a function of temperature for de-excitation transitions to the final 1 01 level of ortho -H 2 S in collisions with (a) para -H 2 ( j 2 = 0) and (b) ortho -H 2 ( j 2 = 1). The initial levels are indicated on the plots. Comparing the plots in Figs 1(a) and 2(a) and the plots in Figs 1(b) and 2(b) , we see that for these lower rotational levels the rate coefficients for collision-induced transitions in ortho -H 2 S are somewhat larger than those in para -H 2 S. 4 DISCUSSION Since H 2 O–H 2 rate coefficients have been used as surrogates for H 2 S–H 2 rate coefficients (Crockett et al. 2014 ; Vastel et al. 2018 ), it is interesting to compare the overall magnitudes of these two sets of rate coefficients. Fig. 3 compares the rate coefficients for de-excitation transitions in ortho -H 2 O and ortho -H 2 S induced by collisions with para -H 2 ( j 2 = 0) and ortho -H 2 ( j 2 = 1) at temperatures of 5 and 100 K. Fig. 4 presents a similar comparison for para -H 2 O and para -H 2 S rate coefficients. The rate coefficients for H 2 S–H 2 collisions are taken from this work, while those for H 2 O–H 2 collisions were obtained from the work by Daniel et al. ( 2011 ). Figure 3. Comparison of de-excitation rate coefficients between the 12 lowest rotational levels of ortho -H 2 O and ortho -H 2 S in collisions with (a) para -H 2 ( j 2 = 0) at 5 K, (b) para -H 2 ( j 2 = 0) at 100 K, (c) ortho -H 2 ( j 2 = 1) at 5 K, and (d) ortho -H 2 ( j 2 = 1) at 100 K. The solid line denotes equal rate coefficients for ortho -H 2 O and ortho -H 2 S, while the dashed lines represent a factor of 2 difference in the rate coefficients. Figure 4. Comparison of de-excitation rate coefficients between the 12 lowest rotational levels of para -H 2 O and para -H 2 S in collisions with (a) para -H 2 ( j 2 = 0) at 5 K, (b) para -H 2 ( j 2 = 0) at 100 K, (c) ortho -H 2 ( j 2 = 1) at 5 K, and (d) ortho -H 2 ( j 2 = 1) at 100 K. The solid line denotes equal rate coefficients for para -H 2 O and para -H 2 S, while the dashed lines represent a factor of 2 difference in the rate coefficients. We see in both Figs 3 and 4 that the transitions in H 2 O with the largest rate coefficients are greater in magnitude than the corresponding transitions in H 2 S, that is to say, the points in upper right-hand corner of these figures are to the right of the solid diagonal line. These differences reflect the different topologies of the H 2 O and H 2 S PES. By contrast, for H 2 S–H 2 transitions with smaller rate coefficients, i.e. those lying on the left-hand side of the plots in these figures, are larger than the rate coefficients for the corresponding transitions in H 2 O–H 2 . It should be noted that the energy spacings in H 2 S are smaller than in H 2 O since the rotational constants for the
Paper: 2018MNRAS.475.1347M (acknowledgments)
ACKNOWLEDGEMENTS We thank the anonymous referee for the valuable comments and suggestions. We thank Andre Izidoro for the discussions concerning our findings during the participation in the ‘Protoplanetary Disks and Planet Formation and Evolution’ programme organized by the Munich Institute for Astro-and Particle Physics (MIAPP). R.A.M. thanks financial support from Fundação de Amparo à Pesquisa do Estado de São Paulo (Grants: 2011/08171-3, 2013/24281-9 and 2016/12113-2).
Paper: 2016A&A...587A..76V (acknowledgments)
We thank the anonymous referee for the critical reading and valuable comments. We thank J. Le Bourlot and B. Godard for the insightful discussions. V.V. acknowledges support from a CNRS-CONICYT scholarship. This research has been partially funded by CONICYT and CNRS, according to the December 11, 2007 agreement. P.H. acknowledges the financial support of the Agence National pour la Recherche through the COSMIS project. This research has received funding from the European Research Council under the European Community’s Seventh Framework Program (FP7/2007-2013 Grant Agreement No. 306483). M.G. and P.L. thank the French Program “Physique Chimie du Milieu Interstellaire” (PCMI). This work was granted access to the HPC resources of MesoPSL financed by the Region Ile de France and the project Equip@Meso (reference ANR-10-EQPX-29-01) of the programme Investissements d’Avenir supervised by the Agence Nationale pour la Recherche
Paper: 2016ApJ...818L..31K (fulltext)
in the formation of van der Waals complexes. The heat of evaporation of a single He atom is about 5 cm −1 (Toennies Vilesov 2004 ). Assuming an extra 7% depletion when Ar gas is replaced with dihydrogen, we can estimate the lower limit of the energy released in the droplets to be about 4200 cm −1 . The real reaction energy is expected to be much higher because not all helium droplets contain both reactants. Additionally, an ejection of the hot reaction products from the helium droplets, as well as a dissipation of the reaction energy by photon emission, are expected (Krasnokutski Huisken 2010b ). This estimate is only consistent with the reaction leading to the formation of the HCH molecule becuase the formation of the CH 2 molecule would lead to the release of only a small amount of energy ∼600 cm −1 (Harding et al. 1993 ). The high rate of this reaction found at ultra-low temperature, assumes no energy barrier in the reaction pathway. Therefore, our experiment is in line with the results of quantum chemical computations that predict the barrierless reaction C + H 2 → HCH (Harding et al. 1993 ). In the case of doping of helium droplets by more than a single dihydrogen molecule, the formed HCH molecule can also further react with a H 2 molecule leading to the formation of methane. However, the quantum chemical computations predict the presence of a barrier for this reaction (Ge et al. 2010 ; Lu et al. 2010 ). Therefore, we expect that this reaction should not proceed inside the helium droplets and that the reactions of carbon atoms are terminated after the formation of HCH molecules. 3.2. Mass Spectrometry In the next step, we applied mass spectrometry for the characterization of products of chemical reactions. In this experiment, larger helium droplets (up to 10 5 He atoms) were used. The droplets were first doped with multiple dihydrogen molecules and, in the second pick-up chamber, one or a few carbon atoms were added. We used two different conditions for the doping with dihydrogen. The helium droplets picked up about 10 and 100 dihydrogen molecules on average at low and high doping conditions, correspondingly. Although, even at the high doping conditions, the C/H ratio is much higher than that present in the ISM, it is low enough to saturate the chemistry with hydrogen (i.e., more than five hydrogen atoms per carbon atom). The doped helium droplets were ionized by electron impact in the head of the TOF mass spectrometer. The collision of a doped helium droplet with an electron, having energy larger than 25 eV, results in the formation of He + . This positive hole migrates within the droplet until it becomes localized either at a dopant or at a helium dimer to form (Scheidemann et al. 1993 ). Therefore, the cations of dopants and helium clusters as well as their complexes are produced. Figure 2 shows the mass spectra of helium droplets doped with dihydrogen alone and together with carbon atoms. Both mass spectra were recorded one after
Paper: 2019ApJ...879...37N (fulltext)
the CH features and C i lines in metal-poor dwarfs with [Fe/H] –3.0 (see Fabbian et al. 2009 ; Jacobson Frebel 2015 ; Placco et al. 2016a ), in what follows we shall assume that this correction is not well-determined below [Fe/H] ≲ −3.0, and make the conservative assumption that A (C i ) 3D,NLTE – A (CH) 3D,LTE = 0.0, for all [Fe/H], and hence A (CH) 3D,NLTE = A (CH) 3D,LTE . The reader should bear in mind that the 3D,NLTE corrections we shall present in what follows are most likely less extreme than would be obtained by adoption of the equation in the previous paragraph. 4. Revised <italic>A</italic>(C) versus [Fe/H] and [C/Fe] versus [Fe/H] Diagrams 4.1. The Yoon et al. (<xref ref-type="bibr" rid="apjab1f84bib113">2016</xref>) Sample of CEMP Stars With the 3D–1D and NLTE–LTE corrections in hand, we now investigate their effect on the distribution of the CEMP-no stars in the ( A (C), [Fe/H]) and ([C/Fe], [Fe/H]) planes. Figure 4 presents the data for the CEMP-no and CEMP-s objects compiled by Yoon et al. ( 2016 ), 7 7 We used only stars in the Yoon et al. ( 2016 ) catalog for which the [C/Fe] values were based on the CH G -band. While this does not affect the CEMP-no stars, it excludes some 15 CEMP-s objects. We also required the presence of [Ba/Fe] to identify membership of the CEMP-no and CEMP-s subclasses, except for stars with [Fe/H] −3.3, where in the absence of detected [Ba/Fe] we assume CEMP-no status. together with those for an additional recently reported eight CEMP-no stars presented in Table 4 , and four stars identified in Table 6 . The effects of the corrections to A (C), [C/Fe], and [Fe/H] are presented in the three rows of the figure. The uppermost row, panels (a) and (b), show data obtained using 1D,LTE for A (C) 1D,LTE and [C/Fe] 1D,LTE versus [Fe/H] 1D,LTE , respectively. Also shown in the top left panel (a) are the the ellipses containing the Groups I, II, and III of Yoon et al. ( 2016 ), together with the “high-carbon band” (horizontal orange line) and the “low-carbon band” (horizontal light blue line) of Spite et al. ( 2013 ), Bonifacio et al. ( 2015 , 2018 ), and Caffau et al. ( 2018 ) (truncated on the right by the [C/Fe] = +1.0 locus). 8 8 We draw the reader’s attention to the fact that the Yoon et al. ( 2016 ) model contains three components, while that of Caffau et al. ( 2018 ) has only two. The question one might ask is: how many components are required to best describe the stellar distribution? We shall discuss this further in Sections 7.1 and 8 . In these panels, (a) and (b), red and gray symbols refer to CEMP-no stars on the one hand, and CEMP-s stars on the other, and in what follows in the middle and bottom rows the symbol for each star will retain the same color as adopted in these uppermost panels. The full and dotted lines in all panels represent the loci of the [C/Fe] > +0.7 divide between C-normal and CEMP stars, and [C/Fe] = 0.0, respectively. Figure 4. Carbon vs. iron abundances for CEMP-no (red symbols)
Paper: 2020ApJ...902...25G (fulltext)
part. The short-wavelength region contains the strong Balmer lines of H γ , H δ , and H ε , while the long-wavelength section only contains H β and weaker lines of He i λλ 4387, 4471, 4713, 4921. The difference in the CCF strength for the two regions suggests that the H Balmer lines are the dominant contributor to the CCF amplitude. In the similar, lower-mass binary, EL CVn, evidence suggests the presence of a thick hydrogen envelope even after extreme mass loss (Maxted et al. 2014 ), and the spectrum of the companion in EL CVn shows strong H Balmer lines with narrow cores (Wang et al. 2020 ), consistent with their appearance in the model CCF template spectra for solar abundances that we use here. Thus, the companion of Regulus probably has a photosphere with a significant hydrogen abundance. 4. Discussion The cross-correlation analysis of the large set of high-S/N spectra from ESPaDOnS/NARVAL reveals a weak signal that has all the properties expected for the spectral signature of the pre-white dwarf companion of Regulus. The CCF peak only appears using model spectra for a hot object, and it attains the strongest and narrowest morphology for a secondary semiamplitude that is consistent with the predicted extreme mass ratio and low mass of the pre-white dwarf remnant. The narrow appearance of the peak agrees with expectations of a low projected rotational velocity for synchronous rotation that is predicted in models of binary evolution. The mass and radius implied by the observations lead to a predicted small gravitational redshift that is consistent with the measured peak velocity. Taken together, these results indicate that the CCF peak marks the first detection of the spectral features of the pre-white dwarf companion of Regulus. The stellar masses we derive from the CCF analysis are similar to those proposed in earlier work (Gies et al. 2008 ). Che et al. ( 2011 ) used the stellar temperature and luminosity and the Yale–Yonsei evolutionary tracks to find a primary mass of M 1 = (4.15 ± 0.06) M ⊙ . This agrees within uncertainties with our result, M 1 = (3.7 ± 1.4) M ⊙ . The pre-white dwarf star has a low mass, , which represents the mass of the surviving core after completion of mass transfer. There is a close relationship between core mass and the properties of the former, Roche-filling red giant, and this implies that the core mass is related to the orbital period of the system. Rappaport et al. ( 2009 ) describe how the relation applies to the Regulus system and how the orbital period leads to a predicted pre-white dwarf mass of M 2 = (0.28 ± 0.05) M ⊙ , and this estimate agrees within uncertainties with the measured mass. The evolution of the stellar core after envelope removal was investigated by Istrate et al. ( 2016 ) using MESA models to study the internal structure and tracks in the plane. We show in Figure 8 one example for a pre-white dwarf with an He core and a mass of 0.305 M ⊙ from the basic set they present, i.e.,
Paper: 2019MNRAS.487.5764T (fulltext)
statistical studies that can identify correlations over longer time-scales, and more accurately measure the scatter between gas reservoirs and black hole growth. It is thus interesting to make predictions for the connection between H i and black hole properties. Fig. 9 shows how black hole properties depend on the H i content of galaxies in simba , specifically the total H i mass M H i (left column), and the H i mass fraction f H i = M H i / M * (right column). The rows show various black hole properties, from top to bottom: M BH , , M BH / M ⋆ , and f Edd . All galaxies are colour coded by sSFR. The running median at z = 0 is shown as the black line with 1σ uncertainties from jackknife resampling. The blue and turqoise lines show the running medians at z = 1, 2 respectively; the individual galaxy points are not shown at those redshifts. Figure 9. Black hole properties as a function of H i content in simba . Each galaxy at z = 0 is colour coded by sSFR and the black line is the running median at z = 0. The blue and cyan lines show the corresponding medians at z = 1, 2, respectively. There is a strong anticorrelation between M BH and f H i , as well as a correlation between f Edd and f H i , both reflecting a tight M BH – M ⋆ relation. The top left panel shows that the black hole mass is essentially uncorrelated with H i mass, and there is little evolution in this relation. More interestingly, at a given H i mass, larger black holes populate more quiescent galaxies. This shows that there is a strong connection between black hole growth and gas removal in galaxy outskirts, likely owing to suppression of cooling to feed the H i reservoir. The most star-forming galaxies primarily have the lowest M BH , but there is also a weaker trend that they have the highest M H i . Hence star formation is enhanced in galaxies that have both small black holes and high gas content. The top right panel shows M BH versus f H i , which displays a strong anticorrelation, reflecting a tight M BH – M ⋆ relation at late times, with the most star-forming galaxies having concurrently the smallest black holes and highest gas fractions. There is modest evolution upwards in this relation with time, such that galaxies at a given black hole mass have higher f H i at earlier times, reflecting the overall increase in gas content in galaxies at earlier epochs (Davé et al. 2019 ). The second row shows the dependence of BHAR on H i content. At z = 0, BHAR shows little correlation with H i mass, but there is an evident correlation at such that black hole accretion is stronger for higher H i masses. At these earlier epochs, the accretion is dominated by the torque-limited mode, which depends on gas fraction. Even though torque-limited accretion is computed within the core of the galaxy while the H i is more diffusely distributed, the overall enhanced gas content appears to drive black hole accretion. By z = 0, in contrast, the emergence of quenched galaxies dominated by Bondi
Paper: 2019ApJS..244....1Z (fulltext)
in Table 2 . For an easy display, only the deviations of the parameters ( a , e , i , ω , Ω, and w ) from the true values (6922.8781 km, 0.00181017, 429715, 2070229, 1169049, and 225215) are listed. Table 2 The Orbit Parameter Spaces Note. “Δ a ,” “Δ e ,” “Δ i ,” “Δ ω ,” “ΔΩ,” and “Δ w ” are the deviations for the semimajor axis, eccentricity, inclination angle, RAAN, argument of perigee, and mean anomaly, respectively. First, the orbit determination method is demonstrated with LE. At each grid point of the parameter space, the “predicted” profile is folded and its significance is obtained. As shown in Figure 3 , the significance of the profile varies significantly with the deviation of the orbital parameters and has a maximum around the true value (zero deviation). Then, the maxima obtained by fitting a Gaussian function curve give the optimal values for the six orbital parameters separately, i.e., each time one parameter is free with the other five fixed, and the results are shown in Table 3 . The process obtaining the best parameters is very similar to that used by Sheikh et al. ( 2007 ) to get the TOAs from the observed pulse profiles (Sheikh 2005 ). Therefore, the uncertainties of these parameter values can be obtained in a similar way as follows:where W is the FWHM, and S/N is the signal-to-noise ratio obtained by the bootstrap method (Diaconis Efron 1983 ). As mentioned by Sheikh (Sheikh 2005 ), the S/N is limited to a maximum of 1000 by Figure 3. Orbit determine result with Insight/LE: χ 2 of the orbit elements of calculations (blue squares) and fitted results (red lines). “Δ a ,” “Δ e ,” “Δ i ,” “Δ ω ,” “ΔΩ,” and “Δ w ” are the deviations for the semimajor axis, eccentricity, inclination angle, RAAN, argument of perigee, and mean anomaly, respectively. Table 3 Best Estimated Value of the Elements and Errors (3 σ ) Then, similar demonstrations are also done for HE and ME, as shown in Table 3 . It shows that the errors for HE are less than those for ME and LE, because it has collected more pulsed photons. By combining all the data from HE, ME, and LE, the errors become smaller. By transferring the orbit elements to the cartesian parameters, we obtain the 3 σ errors of the current position and velocity, which are 3.15 km, 6.61 km, and 4.46 km for x, y, and z ; and 0.0073 km s −1 , 0.0033 km s −1 , and 0.0042 km s −1 for v x , v y , and v z , respectively. In short, the position and velocity are pinpointed within 10 km (3 σ ) and 10 m s −1 (3 σ ). 5. Discussions Conventionally, a minimum of three pulsars are required to get the absolute navigation. Considering the clock time-offset on the spacecraft, another pulsar is needed. Thus, at least four pulsars should be observed simultaneously, which would require at least four detectors, increasing the technical complexity and the cost of the mass, energy, and money. An alternate choice is to observe different pulsars sequentially using a single detector, which would increase the risk of the
Paper: 2018MNRAS.476.3729B (fulltext)
the latter. Fig. 7 shows that if one adopts a T eff lower by 100 K then nearly exact agreement between dynamical and asteroseismic numbers are reached for this star when adopting the correction from Rodrigues et al. ( 2017 ). For KIC8410637 comparisons to very similar stars in the open cluster NGC6819 (Handberg et al. 2017 ) suggests that the observed T eff , measured by Frandsen et al. ( 2013 ), could be too high. There are reasons to suspect that T eff of the binary sample measured by Gaulme et al. ( 2016 ) could be overestimated for some stars, which would cause an overestimate of mass and radius from the asteroseismic scaling relations even if appropriate corrections are applied. The way T eff was measured in that study ignored the continuum contribution from the secondary stars, which makes the spectral lines appear weaker, potentially mimicking a higher temperature. The stars that we re-measured taking into account the light ratio were found cooler, though only by 56, 16, and 12 K. However, the light ratio and therefore the potential overestimate of T eff is larger for the two stars with largest discrepancy between dynamical and asteroseismic measures, KIC5786154A and KIC4663623A, than any of the three giants in our sample. We also note that the spectroscopic log g values reported by Gaulme et al. ( 2016 ) for these two stars are larger than the dynamical and asteroseismic log g measurements by 0.25 and 0.5 dex, suggesting potential issues with the spectroscopic analysis. KIC10001167A can be compared to stars in the globular cluster 47 Tucanae (47 Tuc; NGC104) given that it has similar metallicity ([Fe/H]∼−0.7, compare Gaulme et al. 2016 and Brogaard et al. 2017 ). Since the globular clusters are almost always found to be as old or older than field stars at similar metallicity, KIC10001167A would be expected to have an age equal to or younger than 47 Tuc. A comparison to the turn-off mass of 47 Tuc determined from the eclipsing member V69 (Thompson et al. 2010 ), extrapolated to the giant phase via isochrones (Brogaard et al. 2017 ) would then suggest a mass of for KIC10001167A, about 2σ larger than the measured by Gaulme et al. ( 2016 ). Indeed, if their low mass of is correct, the corresponding age of this star would be larger than the presently established age of the Universe. This problem could be avoided if the star is an AGB star that experienced mass-loss on the RGB. However, we also note that the T eff measured by Gaulme et al. ( 2016 ) for KIC10001167B, the MS star in this binary, is very much (∼6σ) larger than would be expected for their measure of a star on the MS. This suggests that the true mass of KIC10001167B is larger. Since dynamical mass estimates of the two components of a binary system correlate strongly, this also indicates that the mass of KIC10001167A is larger than the dynamical measure. Furthermore, if KIC10001167A is an AGB star then f Δν is ∼1.005 instead of ∼0.957 and thus the seismic mass and radius would be