adsabs/WIESP2022-NER huggingface dataset:
Paper: 2020ApJS..247....4A (fulltext)
can also be seen for physisorption at very low temperature (≤6 K), when the grain can be full of immobile atoms (Katz et al. 1999 ). A corollary of this behavior is that in case the accretion is very efficient, the dust grain could have a large abundance (∼10 −6 ) of a particular species depending upon its desorption energy. For most species, this effect can lead to a change in the gas abundance phase profile. 6. Conclusions We have studied the chemistry of simple molecules on interstellar dust grains bound strongly by chemisorption on graphite. Chemical simulations involving chemisorbed and gas-phase species have been run for both isothermal and warm-up conditions using simplified assumptions. The isothermal models were run between 150 and 400 K at an interval of 25 K, and the warm-up models have a temperature range between 10 and 400 K. For each case, we have run models with different values of adsorption barriers and binding energies to explore the sensitivity of our results to these uncertain chemisorption parameters. Our major conclusions are as follows. 1. The efficiency of chemisorption hinges strongly on the efficiency of adsorption. Models having the lowest adsorption barrier for all species (0.1 eV) were found to produce reasonably abundant chemisorbed species, although the abundance of a chemisorbed species depends both on the adsorption rate and also other processes such as its chemistry. The reduction by a factor of 2 of individual desorption energies calculated by us does not increase the abundances markedly. 2. For species such as CO, which are very efficiently formed in the gas phase at almost all temperatures (10–400 K), the existence of chemisorbed molecules will not have a significant impact on their gas-phase abundances. However, some chemisorbed CO can be locked on grains, leading to a reasonably high abundance of ∼10 −6 , so that it could be detectable using upcoming observational facilities. Species such as C 2 H 2 and NH 3 , which are less efficiently formed in the gas phase compared with CO, may show a change in their gas-phase abundances due to chemisorption. The chemisorbed species can also approach peak abundances near ∼10 −6 and may be detectable. 3. A reasonably large number of chemisorbed species show fractional abundances above 10 −8 under certain conditions. An attempt to detect them using JWST might be worthwhile if chemisorbed species with higher abundances can be detected with reasonable signal to noise. 4. For warm-up models, if the efficiency of adsorption is high, a steplike abundance profile versus time is calculated using the assumption that chemisorption is possible only in the first monolayer, which rules out the formation of more than one layer. 5. Our calculations are based on the scant existing literature and simplified assumptions. In the future, detailed quantum mechanical calculations and experiments will be required to extract essential parameters such as diffusion barriers, sticking
Paper: 2019ApJ...871..218R (fulltext)
strengths where Proca stresses can be significant. As emphasized above, the model considered in this report is not that of dark matter; it may, however, account for the conjectured dark-matter effects in astrophysical settings even beyond galactic rotation curves. It may also be tempting to explore the influence of the negative pressure typical for the Proca electrodynamics on the cosmological-scale phenomena with a possible impact on the dark-energy scenarios. To have a significant effect, one would need the presence of a tangled magnetic field, with a correlation length exceeding the photon Compton length, filling the entire space. The field would have to be nearly force-free to survive on the cosmological timescales. Appendix A Quasi-steady-state Maxwell–Proca Equations A set of Maxwell–Proca equations reads as (see Jackson 1975 ; Goldhaber Nieto 1971 ) Here, and φ are the vector and scalar potentials, ρ and and j are the space charge and the current density satisfying the standard continuity equationNote that here we use the same notation for charge density as we do in the main text for the mass density. This, however, should not cause confusion because the charge density does not appear in any equations in the main text. The electric field is related to and φ byThe potentials satisfy the Lorentz gaugeThis (Lorentz) gauge now becomes the only possible one: Equation ( 54 ) is a direct consequence of Equations ( 50 )–( 52 ). The scalar and vector potentials become measurable quantities. Consider a system where the temporal variation of the “sources” (currents and charges) is slow in the sense that the characteristic time of this variation, , is long compared to the light-propagation time over the spatial scale of the system, L ,(we use the hat overscript to distinguish the characteristic time from temperature). Then, according to Equation ( 48 ), one has the following order-of-magnitude relation:If one associates with L and some velocity that characterizes the motions in the system, then one can say that E is first order in v/ c . Substituting the estimate ( 56 ) into Equation ( 50 ), one sees that the last term on the right is small compared to the first term on the left by a factor of and constitutes a correction of the second order. Neglecting this second-order term, one obtains the quasistatic version of Equation ( 50 ):Likewise, estimating A from Equation ( 49 ) as B × L and using Equation ( 56 ), we find from Equation ( 53 ) an order-of-magnitude estimate . Then, comparing the first and the second terms in Equation ( 54 ), we find that the second term is second-order small and can be dropped, so that Equation ( 54 ) is reduced toTaking the divergence of Equation ( 57 ), one finds that, to the same (first) order, the charge-continuity equation is reduced toThereby we obtain a set of equations describing quasi-steady-state processes in the Maxwell–Proca electrodynamics. The set consists of Equations ( 48 ) and ( 49 ) and Equations ( 58
Paper: 2018MNRAS.479.5325D (fulltext)
channels, thus the results are extremely sensitive to the cross-section of the 27 Al( p , γ) 28 Si reaction: the comparison with the Al measured in E stars seems to require slightly lower rates, but the differences are not sufficiently significant to require a supplementary analysis on this argument. Sodium is the most delicate species entering this discussion, because the surface abundance reflects the equilibria between the production and the destruction mechanisms (see Section 4.3 ), both very sensitive to the temperature, thus exposed to significant variation during the TP phase. This explains the TP evolution of the surface sodium, shown in the right-hand panel of Fig. 3 . As reported in Table 1 , the average sodium of the ejecta, [ Na/Fe ] ∼ 0.5, is similar to the quantity observed in E stars. Despite the vigorous destruction of sodium in the advanced TP phases, the ejecta are enriched in sodium (by a factor ∼3), because during the initial part of the evolution the sodium in the envelope reaches abundances 10–20 times higher than in the matter from which the star formed; this can be clearly seen in the right-hand panel of Fig. 3 . Such a large increase in the surface sodium is partly due to the accumulation of 22 Ne during the SDU (see discussion in Section 4.2 ); an additional motivation is that the sodium enrichment of the envelope takes place during the initial TP phases, when dust production is at the maximum efficiency (region labelled as Mg in the right-hand panel of Fig. 3 ): this favours the release of large quantities of sodium-rich gas into the interstellar medium. 5.3 Can we obtain oxygen-free ejecta? Oxygen deserves a separate discussion. The SAGB yields reported in Table 1 have [ O/Fe ] ∼ −0.2, ∼4 times smaller than FG stars. On the other hand, Fig. 1 above and fig. 3 in Carretta et al. ( 2018 ) show that some stars belonging to group E have extremely low surface oxygen, up to 20 times lower than observed in FG stars; furthermore, some of these measurements are upper limits. We did several tests, by varying the HBB modelling and the description of mass-loss: our conclusion is the impossibility of producing SAGB ejecta with such low oxygen contents. The only possibility is that the stars suffer extremely low mass-loss during the initial TP phases, until most of the oxygen in the envelope is consumed, via p-capture. The rate of mass-loss required to accomplish this task should be below yr −1 , far too small for this high luminosity, expanded stars. In this case, the overall agreement between the yields of SAGB stars and the chemistry of E stars would be substantially worse, because the gas ejected would be practically sodium free and would show magnesium depletion and silicon enrichment far in excess of the quantities observed. On the other hand, it is not possible to reproduce all the oxygen abundances of E stars at the same time, given the large spread observed, amounting to almost 1 dex. Oxygen is very peculiar on
Paper: 2020ApJ...893...96B (fulltext)
ii regions are sensitive to time since the onset of the most recent burst, driving scatter from the average relationships. Henry et al. ( 2000 ) found that chemical evolution models differing only by their assumed star formation efficiencies (SFEs) produced a range of primary N/O plateaus. We illustrate the effect of varying the SFE by over-plotting the Henry et al. ( 2000 ) constant SFR models, where efficiency has been varied by a factor of 25, on our N/O versus O/H data in Figure 14 . For low SFRs, the buildup of oxygen is slow and on the order of the lag time before intermediate-mass stars begin ejecting nitrogen. This allows a high N/O plateau to be established at low oxygen abundances (darkest purple curve). On the other hand, high SFRs early in the SFH form a large number of massive stars that produce greater levels of oxygen ahead of N enrichment, establishing a lower plateau (lightest purple curve) and shifting the entire N/O–O/H trend in Figure 14 to the right toward greater O/H. In between these scenarios, continuous star formation with roughly 250 Myr between bursts will result in N and O increasing in lockstep, dependent on the elemental yields. The coupling of the N/O plateau with galaxy SFH is also reported by cosmological hydrodynamical simulations of individual regions within spatially resolved galaxies (Vincenzo Kobayashi 2018 ). In these simulations, asymptotic giant branch (AGB) stars contribute significant N at low O/H, but the exact value of the primary N/O plateau will vary from galaxy to galaxy according to the relative contributions from SNe and AGB stars, as determined by their galaxy formation time and SFH. On the left-hand side of Figure 14 , we extend the highest N/O plateau from NGC 5194 (green) and the lowest N/O plateau from NGC 5457 (yellow). Based on the above discussion, for the low N/O plateau of NGC 5457, we can put forth an SFH scenario in which the SFE was high early in the galaxy’s evolution, allowing oxygen to build up from many bursts of star formation before nitrogen was returned from longer-lived intermediate-mass stars. Due to the higher level of nucleosynthetic products from massive stars, contributions from secondary nitrogen production may dominate over primary nitrogen production at relatively low O/H and S/H values. On the other hand, the high N/O plateau of NGC 5194 could be due to an SFH in which low SFE at early times allows nitrogen production, although delayed, to keep pace with oxygen and sulfur production and enrich the ISM. Here, we assume low SFE to mean either constant, low SFRs or long quiescent periods between bursts. In this scenario, primary nitrogen production is the dominant mechanism until the galaxy reaches relatively high O/H. Note, however, that this is a very simplistic model where N/O is changing monotonically, in a hierarchical galaxy-building scenario that may not be true. In summary, the primary N/O plateau sensitively probes the SFH of a galaxy, rather than being set by
Paper: 2018MNRAS.473.2465B (fulltext)
The initial sample contains 777 spectroscopically unobserved SuperWASP/ ROSAT objects, with rotation periods ≲ 5 d, declinations >−20°, 8 V 15 and spectral-types FGK, corresponding to 0.7 V − K s 3.8. Targets in the first list are selected from the parent sample. A second list comprised of a subset of 160 targets whose proper motions and positions are compatible with membership of the Pisces MG. B and V photometry for all targets was sourced from the catalogue that provides the smallest errors: APASS DR9 (Henden et al. 2016 ), UCAC4 (Zacharias et al. 2013 ) or NOMAD (Zacharias et al. 2005 ). All K s photometry is from 2-Micron All-Sky Survey (2MASS; Cutri et al. 2003 ). 2.1 Selecting potentially young stars in superWASP We learned from BJM15 that the majority of the raw light curves of young stars exhibited irregular features as well as sinusoidal behaviour, indicative of star-spot modulation and/or chromospheric plages – both prognostic of stellar youth. We not only assigned priority to any objects whose light curves displayed any of these features but also demanded clean light curves and consistent rotation period measurements over several seasons of observations. We also tried to screen out, through manual inspection of the light curves, targets that looked like eclipsing binaries. The reader is referred to section 2 of BJM15 for a description of the SuperWASP project and post-processing techniques used to measure rotation periods, which are identical to those carried out in this work. Table 1 shows the details of the period measurements for the sample of likely-young and single objects in our sample, which are defined in Section 5.7 . Fig. 1 displays rotation periods as a function of V − K s colour for our entire observed sample. Details for the other observed targets are provided online. We obtained spectroscopy for 122 stars, 103 selected from the criteria described here, and a further 19 satisfy these criteria, but have proper motions consistent with membership of the Pisces MG (see the next section). Figure 1. Rotation periods and V − K s colour for the entire sample observed in this paper. Red squares represent 17 objects that were identified to be likely-young stars and were considered likely-single (referred to as the ‘likely-young sample’, see Section 5.7 ). These objects are labelled from 1 to 17 and their corresponding SuperWASP names are provided in Table 1 . Any objects that had two or more separate radial velocity measurements whose mean values varied by more than 2 error bars from each other were assumed to be short-period binaries (see Section 4.2 ) and are denoted here by open green circles. All other objects are shown with filled green diamonds. This symbol scheme is used throughout this work, unless otherwise stated. The various lines represent the I- and C-sequence gyrochrones proposed by Barnes ( 2003 ) and calibrated by Mamajek Hillenbrand ( 2008 ), plotted at a range of ages. Table 1. Measured periods for the
Paper: 2020MNRAS.498.2196S (fulltext)
kinematics study of Clarkson et al. ( 2008 ). Awiphan et al. ( 2016 ) presented a more recent version of the model (variant BGM1307) and compared it with a large ensemble of 470 events from the MOA-2 survey (Sumi et al. 2013 ). Awiphan et al. ( 2016 ) concluded that there was a reasonable level of agreement once allowance was made for brown dwarfs that are missing from the BGM. However, the model appeared to underpredict the observed rate from the Galactic bulge significantly. Sumi Penny ( 2016 ) subsequently showed that at least part of the cause of this discrepancy was undercounting of source stars within the MOA-2 analysis and correcting for this resulted in much better agreement with the model. After investigation of a number of different BGM variants, we have decided to use BGM1307 as the basis for MaBμlS-2. The BGM separates the Galaxy into four components; the thin disc, thick disc, bar and halo, each with their own stellar initial mass functions (IMFs), star formation rates (SFRs), kinematics and ages. Interstellar extinction is included for UBVRIJHK photometric bands using a 3D extinction model from Marshall et al. ( 2006 ). The solar position relative to the Galactic Centre and plane in the BGM is taken to be R 0 = 8 kpc and z 0 = 15 pc, respectively. The following description of the Galactic components quotes lower limits on their respective stellar mass functions, but an upper limit is not included. While interesting lens candidates such as intermediate-mass black holes and other compact objects are represented in the high-mass regime, the BGM does not include them. Due to the steepness of the mass functions at high mass, these objects also do not contribute significantly to the overall microlensing parameters outlined in Section 3.1 . 2.1 The thin disc The thin-disc component is subdivided into seven age groups, each with their own luminosity and effective temperature distributions. The mass density is modelled by the difference of two exponentials in cylindrical polar coordinates (r, z) ; one representing the body of the disc with scalelength R d = 2170 pc and the other representing a central hole with scalelength R h = 1330 pc (Robin et al. 2012b ). It is given by where ρ 0 is the component normalization density, and ε is the axial ratio of the ellipsoid. The SFR is taken to be constant across the whole disc. The IMF is given by a broken power law, which is the same across all age groups : Each age group of the thin disc has its own velocity dispersion, to mimic the secular evolution, and follows the relation from Gomez et al. ( 1997 ); these average to km s −1 , with a maximum of km s −1 for the oldest component. The rotational velocity of the local standard of rest (LSR) is 240 km s −1 , with the rotation curve taken from Caldwell Ostriker ( 1981 ). The LSR is defined as the mean velocity of material in the solar neighbourhood, centred on the Sun within 100 pc. The peculiar motion of the Sun relative to the LSR is assumed to
Paper: 2019MNRAS.489.5494T (acknowledgments)
ACKNOWLEDGEMENTS DFT and FCZ acknowledge support from the Spanish Ministry of Economy, Industry and Competitiveness grants PGC2018-095512-B-I00, SGR2017-1383, and AYA2017- 92402-EXP. DV acknowledges support from the Spanish Ministry of Economy, Industry and Competitiveness grants AYA2016-80289-P and AYA2017-82089-ERC (AEI/FEDER, UE). FCZ is supported by a Juan de la Cierva fellowship. DFT acknowledges the participants of the Workshop on Magnetospheres of Neutron Stars and Black Holes (June 2019, Goddard Space Flight Center) for comments.
Paper: 2018ApJ...861...46H (acknowledgments)
S.M.H. acknowledges support from Ramanujan Fellowship SB/S2/RJN-73 and the Max-Planck Partner group program. S.M.H. is grateful to Martin Woodard, with whom many useful conversations were had, for helping resolve several technical issues. This research was supported in part by the International Centre for Theoretical Sciences (ICTS) during a visit for participating in the program—Turbulence from Angstroms to Light Years (ICTS/Prog-taly2018/01).
Paper: 2015ApJ...814...89W (acknowledgments)
This paper is based on data collected at the Subaru Telescope, which is operated by the National Astronomical Observatory of Japan. We thank the anonymous referee for helpful suggestions and Takashi Hattori for taking the data for us. This research was supported in part by the National Natural Science Foundation of China (11373055) and the Strategic Priority Research Program “The Emergence of Cosmological Structures” of the Chinese Academy of Sciences (Grant No. XDB09000000). C.W. acknowledeges the support from the National Natural Science Foundation of China (11273029). Facility : Subaru (FOCAS)
Paper: 2020ApJ...900...96A (fulltext)
efficiency of sublimation in this context, compared with the incidence of “dry rain” in the Earth’s atmosphere, can be ascribed to a combination of two factors: the relatively high silicate vapor mass fraction (typically X sil ∼ (0.3–1) X H when refluxing of pebble debris from the envelope begins to balance pebble accretion) and an infall time that is a factor ∼10 3 times longer than for raindrops. Of course, much larger planetesimals that survive transport through the envelope may still inject solid/liquid silicates into the core. 3. Envelope Expansion in Response to Planetesimal Heating and Dust Loading: Semianalytic Treatment As a prelude to the numerical investigation described in Sections 4 and 5 , we now present a semianalytic model of an envelope with a low dust density and outer radiative structure. This model connects smoothly to a uniform medium of a fixed temperature T disk and gas density ρ g ,disk . As the dust abundance rises, the inner convective layer grows in size, and eventually overwhelms the Bondi radius. Assuming that the dust abundance in the outer radiative layer is tied directly to the solid accretion rate, we can work out the critical value of above which convection extends beyond the Bondi radius, some refluxing of solids will occur, and the self-consistency of the model breaks down. We next quantify the reduction in dust size due to convective stirring, and demonstrate how a convective layer that traps accreted solids may expand in response to deep heating by large planetesimals. 3.1. Two-layer Convective-radiative Envelope Even a small amount of embedded dust acts as a powerful coolant, in the sense that any excess of gas temperature over ambient radiation temperature is rapidly converted into thermal radiation via collisions between gas molecules and grains (e.g., Chiang Goldreich 1997 ). The presence of an ambient blackbody radiation field of temperature T disk is therefore assumed, with an internal energy amounting to a small fraction of the ambient gas energy density. The temperature is buffered near T disk in the outer bound envelope, around the Bondi radius, and grows in the inner, optically thick envelope. In a state of spherical hydrostatic and radiative equilibrium at local temperature T g and uniform luminosity ,The opacity is per unit mass of dust. We first consider the case of uniform dust-mass fraction X d (although the equations presented here do not depend on this assumption). We rescale radius to R B , temperature to T disk , and gas density to ρ g ,disk . Then the preceding equations can be writtenwhereand the luminosity is expressed in terms of a core-accretion time,Taking cm 2 g −1 (see Appendix A ), the characteristic optical depth at the Bondi radius isand The solution to Equation ( 25 ) is shown in Figure 2 for β = 0 and a range of values of the parameter ε rad (10 −4 , 10 −3.5 , 10 −3 , ... 10 0.5 , 10). When ε rad ≪ 1, the temperature remains flat some distance inside R B , whereas the