adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016ApJ...829...34S (fulltext)
this correlation is presently unclear. Figure 8. Periastron distance (top panel; ) and the tidal force acting on the primary star at periastron divided by the star’s surface gravity (bottom panel; see Equation ( 4 )), both as a function of the full amplitude of the photometric HB signal at periastron. In both panels the markers’ radius is linear in the primary star’s (See Table 1 ). 4.5. Higher Multiplicity Systems Many of the HB stars examined here may be members of triple or higher multiplicity systems. Several works (e.g., Duquennoy Mayor 1991 ; Meibom Mathieu 2005 ; Tokovinin et al. 2006 ; Raghavan et al. 2010 ) have found that the fraction of higher multiplicity systems among short period and highly eccentric binaries is very high, exceeding 90% for binaries with days (Tokovinin et al. 2006 ). Indeed, tertiary bodies may excite the orbital eccentricity of HB progenitors via Kozai–Lidov oscillations, producing close periastron passages and allowing them to be detected as HB systems. Quadruple systems composed of two binaries may also be common among HB systems, and these types of systems have been predicted by Pejcha et al. ( 2013 ), who found that HB systems formation is greatly enhanced in quadruple systems relative to triple and binary systems. Despite the probable multiplicity of many HB systems, we find it unlikely that many of our Keplerian RV solutions have been greatly affected by orbital motion induced by a third body. A putative third body exciting the orbital eccentricity would likely have an orbital period, P 3 , more than several times longer than the HB period (depending on its mass, inclination, and eccentricity) in order to allow long term dynamical stability (Kiseleva et al. 1994 ). If a typical HB progenitor system was born with an orbital period days before having its orbital eccentricity excited, we expect for most third bodies, with much larger values of P 3 in most cases. Although the orbital RV amplitude induced by such a third body can be several km s −1 , we expect only a small RV change over the ∼100 days baseline of our observations, and therefore this motion can be safely included as a constant RV offset. We are currently pursuing a follow-up survey on many of these HB systems that will reveal whether any of them have tertiary companions with orbital periods . 4.6. Future Prospects The orbits measured here comprise a relatively large sample obtained with a single instrument, and they can facilitate several follow-up scientific studies in addition to that of tidal circularization and the shape of the e – P distribution (Section 4.3 ). Other examples include: 1. Stars in binary systems with highly eccentric orbits are expected to rotate near a “pseudo-synchronous” rotation period, where the stellar rotation is synchronized with the orbital motion close to periastron which is faster than the mean orbital motion along the entire orbit. The pseudo-synchronous rotation rate depends on eccentricity and orbital period
Paper: 2019ApJ...871..167K (acknowledgments)
A.F.K. acknowledges support from University of Maryland GPHI Task-132, HST GO 13323, NASA Exoplanet Science Institute (NASA/Keck telescope time), an appointment to the NASA Postdoctoral Program at the NASA’s Goddard Space Flight Center, administered by Universities Space Research Association (previously by the Oak Ridge Associated Universities) under contract with NASA. A.F.K. thanks Joel Allred and Mats Carlsson for many helpful discussions about RADYN, Han Uitenbroek for helpful discussions about the RH code, and Pier-Emmanuel Tremblay for the hydrogen broadening profiles. A.F.K. thanks James Davenport for observations from the ARCSAT 0.5 m and for providing a flare rate for GJ 1243, Nicola Gentile Fusillo for observations from the Isaac Newton Telescope, Mihalis Mathioudakis for helpful discussions about the ULTRACAM data, Lyndsay Fletcher for helpful discussions about the interpretation of the HST data, and Lucianne Walkowicz for helpful feedback on the initial HST proposal. We thank R. O. Parke Loyd and the STScI helpdesk for helpful feedback on the jump in the dispersion solution in the COS spectra. A.F.K. also thanks the Keck Observatory support astronomer Luca Rizzi for assistance with the observations. IRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy (AURA) under a cooperative agreement with the National Science Foundation. Some of the data presented in this paper were obtained from the Mikulski Archive for Space Telescopes (MAST). STScI is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. Support for this work was provided by NASA through grant No. Guest Observer 13323 from the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS 5-26555. Based on observations obtained with the Apache Point Observatory 3.5 m telescope, which is owned and operated by the Astrophysical Research Consortium. The 2.3 m Aristarchos telescope is operated on Helmos Observatory by the Institute for Astronomy, Astrophysics, Space Applications and Remote Sensing of the National Observatory of Athens. This work has made use of data from the European Space Agency (ESA) mission Gaia ( https://www.cosmos.esa.int/gaia ), processed by the Gaia Data Processing and Analysis Consortium (DPAC https://www.cosmos.esa.int/web/gaia/dpac/consortium ). Funding for the DPAC has been provided by national institutions, in particular the institutions participating in the Gaia Multilateral Agreement.
Paper: 2016ApJ...829...86S (fulltext)
tidal forces from the Milky Way, or any other physical process 9 9 For example, relaxation cannot be responsible for the spatial distribution of the different stellar populations in Carina since the relaxation times in dwarf galaxies are much larger than the age of the universe (e.g., Freeman 2008 ). , were not strong enough to move a significant fraction of these stars to the outer region. Otherwise, a large number of intermediate-age stars would be present in the outskirts of Carina, and thus we would not observe such a strong break in the SFH and MDF of Carina between the intermediate and outer regions. Since the radial trends in Carina do not seem to be governed by the tidal influence of the Milky Way, we suggest that this hints that the SFH is not dominated by tidal forces either. In this line, the fact that we find two episodes of star formation in Carina instead of three, as has often been claimed (Hurley-Keller et al. 1998 ; Rizzi et al. 2003 ; Pilkington Gibson 2012 ; McMonigal et al. 2014 ), makes it harder to explain the star-formation episodes in Carina as being the result of periodic close passages to the Milky Way (as done in Pasetto et al. 2011 ), even more so since proper motions suggest that the orbital period of Carina is close to 2 Gyr (Piatek et al. 2003 ; Pasetto et al. 2011 ). Another mechanism proposed to explain the SFH of Carina is gas accretion from the Milky Way (e.g., Lemasle et al. 2012 ; de Boer et al. 2014 ). In this view, the first episode would have ended because of gas depletion, and subsequent star formation would have been stopped for a few gigayears until the dwarf galaxy renewed its gas content thanks to an inflow from the Milky Way. Even though this is a plausible scenario, the in-falling gas should have had the same metallicity as the stars from the end of the first episode to reproduce the age–metallicity relation in Carina. A simpler way to explain why the metallicity of the stars at the end of the first episode coincides with that of the stars at the beginning of the second episode is to assume that the gas that formed the second episode in Carina was enriched within the same galaxy by the first generation of stars. Internal evolution would also be consistent with the mean metallicities and ages at different Carina regions, which according to what we see in Figure 9 are roughly constant throughout the extent of the galaxy. This would not necessarily be the case if the gas that formed those stars had been accreted at different moments and each region had different levels of influence from each inflow. For all these reasons, we conclude that the SFH of Carina must be dominated by internal evolution. Our interpretation is that, after the first episode formed the majority of the stars in Carina, the physical conditions of the gas made the galaxy unable to form stars for a few gigayears, until it contracted to activate star formation again. We claim that the temporal gap in star formation in Carina might
Paper: 2019MNRAS.482..732C (fulltext)
of the disc will be In our model, the relative change in disc mass and accretion rate in an encounter is determined by γ and the truncation radius relative to the disc characteristic radius just before the encounter, . If the disc is truncated, its viscous time-scale decreases and the disc starts evolving faster. We ignore the orientation of the discs, so the equation for truncation radius is the average truncation radius over all disc inclinations. 2.4 Numerical implementation We use the fourth-order Hermite N -body code ph4 , which is incorporated into the AMUSE framework. The softening length is and the time-step parameter is η = 10 −2 . The simulations start in virial equilibrium and last for 2 Myr. Collisional radii for each star in the cluster are defined depending on the stellar mass and the theoretical size of the disc evolving in isolation, given by equation ( 1 ). The initial collisional radius for each star is given by the distance at which the most massive star in the cluster can truncate its disc. The collisional radii of all stars are updated every 2000 yr to account for the viscous evolution of the discs. Two stars are considered to be in an encounter if the distance between them is less than the sum of their collisional radii. The encounter distance between the stars is then computed by analytically solving the two-body problem. If the encounter is strong enough as to truncate the discs, disc parameters of both stars are updated as described in Section 2.3 . After each encounter, the collisional radii of both stars are set to 0.49 r enc in order to prevent the N -body code from detecting the same encounter multiple times during the same 2000 yr window. 3 RESULTS 3.1 Initial conditions 3.1.1 Cluster properties Simulations were run for clusters with 1500 stars. The stellar masses are randomly sampled from a Kroupa initial mass distribution (Kroupa 2001 ) with lower mass limit 0.1M ⊙ and upper mass limit 100M ⊙ . The stars are initially distributed in a Plummer sphere (Plummer 1911 ) with Plummer radius 0.5 parsec. All stars are assumed to be single and coeval. All simulations were carried out for . We define three different gas scenarios for the simulations: No gas Gas presence Gas expulsion starting at Each simulation was run for two values of the turbulence parameter: α = 10 −2 and α = 5 × 10 −3 . 3.1.2 Disc properties Based on the observations of low-mass stars carried out by Isella et al. ( 2009 ) and the estimations obtained by Hartmann et al. ( 1998 ), the initial characteristic radius R c of the circumstellar discs was chosen to be with . Circumstellar discs with masses larger than 10 per cent of the mass of their star are the most likely to be gravitationally unstable (Kratter Lodato 2016 ). According to hydrodynamical simulations of collapsing gas clouds, the disc-to-star mass ratios of embedded protostars are large enough for gravitational instabilities to occur (Vorobyov 2011 ), but these instabilities lead to
Paper: 2016MNRAS.462.2878P (fulltext)
from the MPA-JHU DR7 catalogue, and we obtained MBH only for galaxies classified as ellipticals (see Section 3 ). The values are given in Table 8 . MLLINERs cover the range between log (MBH/M ⊙ ) = 7.03–8.57 with a median value of log (MBH/M ⊙ ) = 7.45, while LLINERs show MBH in the range log (MBH/M ⊙ ) = 6.24–8.54 and median value of log (MBH/M ⊙ ) = 8.04. Interestingly, our MLLINERs do not contain the most massive BHs in their centres. It could be surprising that MLLINERs, having on average higher SFRs than LLINERs, show in general lower stellar masses. Different works, both observational (Kauffmann et al. 2003b ; Mateus et al. 2006 ; Leauthaud et al. 2012 ; Pérez et al. 2013 ) and numerical (Shankar et al. 2006 ; Behroozi et al. 2012 ), revealed a stellar mass of ∼6 × 10 10 M ⊙ as critical for the growth rate of stellar populations. In particular, in Pérez et al. ( 2013 ) by studying a 3D spectroscopic sample of 105 local galaxies, the authors found that in galaxies more massive than 5 × 10 10 M ⊙ the inner regions ( 0.5 R 50 ) grew as much as 50–100 per cent faster than in the lower mass galaxies. They found that the peak of relative growth rates of inner and outer galaxy regions correspond to the stellar mass of 6–7 × 10 10 M ⊙ (see their fig. 5), while for lower and higher masses the growth rate decreases and therefore SFRs (LSF) as well. The median stellar mass of our MLLINERs (6.58 × 10 10 M ⊙ ) corresponds perfectly to this region, while for most LLINERs their stellar masses are already higher and correspond to lower values of the relative growth rate (lower LSF). This explains why MLLINERs having in average lower stellar masses in comparison to LLINERs, have higher LSF. In addition, we studied the stellar mass distributions of MLLINERs and LLINERs for the three morphological groups. Table 9 shows the median stellar masses for different morphological types. As can be seen, of three morphological types the highest difference was obtained for early-type galaxies. These galaxies represent a significant fraction of MLLINERs (40 per cent, see Fig. 8 ) and their median mass corresponds exactly to the highest relative growth rate of stellar populations (according to Pérez et al. 2013 ), which then could explain their high SFR values. This is not the case for early-type LLINERs that are characterized by higher stellar masses (lower growth rates) and lower SFRs. Table 9. Median total stellar masses of MLLINERs and LLINERs in relation to morphology (given as logarithm and in M ⊙ ). 5.1.2 Extinction We found that our MLLINERs can be hosted by galaxies with a wide range of extinctions. When using the Av measurements based on the Hα and Hβ emission lines, we find that most of them reside in galaxies with high extinctions. The median Av is 1.65 mag, covering the range 0.49–3.46 mag (see Fig. 7 , top plot). For comparison, the bottom plot in Fig. 7 shows the Av distributions of LLINERs and MLLINERs when taking into account the SDSS MPA-JHU DR7 3
Paper: 2015A&A...584A..60C (acknowledgments)
The authors acknowledge research funding by Deutsche Forschungsgemeinschaft (DFG) under grant SFB 963/1, project A5.
Paper: 2015MNRAS.446.3885B (acknowledgments)
The authors would like to thank the anonymous referee for helpful comments and suggestions. The scientific results reported in this article are based on data obtained from the Chandra Data Archive. We also made use of the ROSAT data archive of the Max Planck Institute für Extraterrestrische Physik (MPE) at Garching, Germany and NRAO VLA Sky Survey archives.
Paper: 2019ApJ...885...22T (acknowledgments)
This work was supported by the grant of Universiti Kebangsaan Malaysia (GUP-2018-086) and Grant-in-Aid for Scientific Research (C) 17K05673 from the Japan Society for the Promotion of Science (JSPS). W.L.T. would like to thank Takuma Nakamura and Bo-Jhou Wang for useful discussions. Special thanks to the dedicated efforts of the entire MMS mission team for data access. MMS data can be downloaded through the NASA CDAWeb at http://cdaweb.gsfc.nasa.gov/ .
Paper: 2020MNRAS.492..283H (fulltext)
environments. Our results will aid in the preparation of multicomponent ices containing acetaldehyde by allowing for calibrations and determinations of that compound in an ice sample. We also envision applications of our data to the measurement of reaction yields in experiments that might produce acetaldehyde in polar ices. Our IR intensities also can be used to estimate band strengths in spectral regions beyond the mid-IR used in the present work. See Gerakines et al. ( 2005 ) for an example. Astrochemical results based on IR band strengths published before those presented here can now be updated with our new measurements. For example, and as already described, the reference band strength of Schutte et al. ( 1999 ) for amorphous acetaldehyde's CH 3 deformation at 1350 cm −1 is A ′ = 1.5 × 10 −18 cm molecule −1 , which now is seen to be in error by nearly 80 per cent. Molecular abundances or column densities based on that older A ′ will be in error by that much, although published relative band strengths might not change substantially. Finally, we again point out that few aldehydes have been examined with the methods used here. As an example, propionaldehyde (propanal) is an interstellar molecule, and at least two laboratory papers recently have been published involving it (Jonusas, Guillemin Krim 2017 ; Qasim et al. 2019 ), but no band strength measurements are available for the solid compound. The situation is no better for H 2 CO (formaldehyde). Many laboratory papers have been published involving H 2 CO, but only estimates and indirect measurements are available for its IR intensities. Unfortunately, their accuracy is unknown. 5 SUMMARY AND CONCLUSIONS Here we have reported the first direct measurements of band intensities of solid acetaldehyde, removing the need to rely on either calculated gas-phase results or values that cannot be traced to their original source. We also have presented mid-IR spectra of the two crystalline polymorphs previously identified by neutron-diffraction studies, one stable under our conditions and one metastable. The IR spectrum of H 13 C(O) 13 CH 3 has been recorded and used to identify ketene as a radiolytic decomposition product of solid acetaldehyde for the first time. APPENDIX A The procedure used to determine absorption coefficients ( α ′) and band strengths ( A ′) is the same as in our earlier papers, and is essentially the method of Hollenberg Dows ( 1961 ). Briefly, IR spectra were recorded of ice samples having a range of thicknesses. The absorbance at each peak of interest was measured and bands of interest were integrated. Multiplication of each absorbance or band area by ln(10) ≈ 2.303 was then done to convert from a common to a natural logarithmic scale. Representative plots of the results are shown in Fig. A1 . The slope of the upper graph is α ′ and the slope of the lower one is ( A ′ ρN A / M ), where M = molar mass = 44.05 g mol −1 , N A = 6.022 × 10 23 molecules mol −1 , and density ( ρ
Paper: 2020ApJ...894...54D (fulltext)
(Linder 2008 ). 3.4. Growing <italic>ν</italic> Mass Growing neutrino mass models, wherein the mass of the neutrino ( m ν ) increases with time and stops the dynamical evolution of the dark-energy scalar field are invoked to solve the cosmological coincidence problem, i.e., the problem that the present-day matter density and density of Λ are similar order of magnitude despite their different dependence on the scale factor (Fardon et al. 2004 ; Wetterich 2007 ). The combined dark sector (scalar field plus neutrinos) energy density (where a = 1/(1 + z )) is given by where Ω ds = 1 − Ω M is the present-day dark-energy density. The scale factor at which the neutrinos become nonrelativistic, is given by preserving continuity between the early and late-time terms (i.e., setting Equations ( 14 ) and ( 15 ) at a = a t ). The two free parameters are the early dark-energy density Ω e and the neutrino density Ω ν . The normalized Hubble parameter for this model is given by 3.5. Dark-energy Transitions at Low Redshift Mortonson et al. ( 2009 ) proposes models with large fluctuations in the dark-energy equation of state at low redshifts, typically z ≲ 0.02, that induce step-like transitions in the H ( z ). Such changes are hidden from constraints coming from Hubble flow SNe Ia (as even the lowest redshift SNe can typically be at z > z t ), but calibrators offer an additional restriction. Step-like responses in H ( z ) also evade several model-independent constraints, as they often assume smoothness of H ( z ). Phenomenologically, we can write the first Friedmann equation in these models aswhereIn this definition, z t , and δ are the position, width, and amplitude of the transition, respectively. In particular, δ = 0 corresponds to the standard ΛCDM . Hence, this model has three free parameters at a given fixed z t : , and δ . The observed Hubble constant is . In this work, we either assume z t = 0.02 or z t = 0.1; the former case motivated by Mortonson et al. ( 2009 ) and the latter case motivated so we can use all the low- z SNe at z ≈ 0.1. The adopted priors for and are equivalent to the priors shown on Table 1 for H 0 and Ω M . The difference in notation highlights the fact that and cannot be interpreted as the observed Hubble constant and the dark-matter density at redshift zero in these models (see Mortonson et al. 2009 ). 3.6. Cosmographic Expansion Along with the different dark-energy models described above, we also look into a more model-independent method, by expanding the expression for H ( z ) as a Taylor series in a cosmographic approach (for e.g., Arendse et al. 2019a , 2019b ; Feeney et al. 2019 ; Macaulay et al. 2019 ). Expanding H ( z ), we getwhere , 2 and . 4. Results In this section we present the results of fitting the different dark-energy models described in Section 4.1 to the combined calibrator and Hubble flow SNe Ia. We also discuss the impact of the systematics covariance matrix on the inferred value of H 0 in Section 4.2 .