adsabs/WIESP2022-NER huggingface dataset:
Paper: 2021MNRAS.500.5117R (fulltext)
L183 (L134N), Taurus, and IC 5146). In the discussion that follows, we must recognize that the observed (i.e. line of sight) extinction to a source is not necessarily the same as the ‘local’ extinction, that is, to say the extinction of the source relative to the interstellar radiation field, so there will always be some uncertainty/natural scatter in the values of A th . However, the near-universality of these dependencies implies that the relevant processes must be largely independent of the dust properties; such as the size distribution and the dust grain morphology and porosity, etc. In the following section, we describe the nature of interstellar ices and the criterion for ice mantle growth. In Section 3 , we describe the various microscopic mechanisms that determine A th , and in Section 4 we compare the relative importances of these processes. Section 5 describes our chemical model, and the results are presented in Section 6 . Possible alternative causes of the A th phenomenon are discussed in Section 7 . Our conclusions and discussion are presented in Section 8 . 2 THE REQUIREMENT FOR ICE MANTLE GROWTH Interstellar ices are predominantly composed of H 2 O, CO, and CO 2 , with CO and CO 2 having abundances of ∼27–32 per cent and ∼19 per cent, respectively, relative to H 2 O in Taurus and low-mass young stellar objects (YSOs) (Chiar et al. 2011 , table 4). CO, being more volatile, is typically less abundant in high mass YSOs. Other species (such as CH 3 OH, NH 3 , OCN − , and other organics/COMs) typically account for ≲10 per cent of the ice. The relative abundances of the major ice components are fairly similar in different environments, although real differences exist and there are often significant variations in the abundances of the complex organics. Whilst there is evidence for significant source to source variations in the relative abundances of H 2 O, CO, and CO 2 , the total contribution of O-bearing ice molecules to the oxygen budget follows similar trends in different souces (e.g. Whittet et al. 2009 ). The chemistry of the main components is straightforward; CO ice is believed to be formed (primarily) due to the freeze-out of CO from the gas phase. By contrast, H 2 O and CO 2 are probably formed by surface reactions. At least four distinct H 2 O formation mechanisms are possible (van Dishoeck et al. 2013 ), but in the relatively low density/high hydrogen atom abundance environments that we are considering in this study, the main channel is simple hydrogenation: The latter reaction being in competition with As gas-phase chemistry is not believed to be important, this means that the formation rate of H 2 O ice is mainly determined by the rate at which oxygen atoms stick to grains and are subsequently hydrogenated to H 2 O molecules. If these are the only relevant reactions, then it also follows that the abundances of H 2 O and CO 2 ice should be approximately correlated, which is indeed seen to be the case (e.g. Noble et al.
Paper: 2016A&A...587A..91B (fulltext)
region. Several studies ( Jacq et al. 1990 , 1999 ; Lubowich et al. 2000 ; Polehampton et al. 2002 ) suggested that deuterium may be up to 10 times less abundant in the Galactic center region than in the local interstellar medium owing to stellar processing. The levels of COM deuteration found in Sgr B2(N2) are indeed systematically lower (by a factor of a few up to a factor 10) than those reported in the Orion-KL Hot Core or Compact Ridge. While a chemical or evolutionary origin for this difference cannot be excluded without further modeling, the low deuteration level of complex organic molecules in Sgr B2(N2) may result from deuterium being less abundant in the Galactic center region by maybe up to a factor ten. 7. Conclusions We searched for deuterated complex organic molecules in a complete 3-mm interferometric line survey performed toward the hot core Sgr B2(N2) with ALMA. We report the secure detection of CHDCN and tentative detections of CHDOH, CHDCHCN in its out-of-plane conformation, the chiral molecule CHCHDCN, and DCN. We also derive column density upper limits for CHOD, CHDCHCN in its in-plane conformation, the three deuterated isotopologues of vinyl cyanide, the four deuterated species of ethanol, and CHDOCHO in its in-plane and out-of-plane conformations. The detections are characterized by compact emission consistent with a hot-core origin. We obtain a deuterium fractionation level of 0.4% for CHDCN and values between 0.05% and 0.12% for the other (tentatively detected) species. Stringent upper limits are derived for the in-plane conformer of CHDCHCN () and CHOD (). The deuteration upper limits for the other nondetected deuterated species lie in the range 0.5–1.8%. The following conclusions arise from this analysis: 1. Ethyl cyanide is less deuterated than methyl cyanide by at least a factor five. This may be due to their different formation pathways (grain surface versus gas phase). 2. The relative abundances of the three deuterated species of ethyl cyanide are most likely consistent with a statistical distribution. 3. The [CHDOH]/[CHOD] abundance ratio in Sgr B2(N2) is higher than 1.8. It may still be consistent with the value obtained in Orion KL but a value closer to (or, as in low-mass Class 0 protostars, even higher than) the statistical value cannot be excluded. 4. Except for methyl cyanide, the levels of deuterium fractionation measured toward the hot core Sgr B2(N2) lie at least a factor four below the predictions of current astrochemical models that were designed for lower mass objects (hot corinos in the Class 0 phase). 5. The levels of deuterium fractionation derived for complex organic molecules in Sgr B2(N2) are lower by a factor of a few up to a factor ten than in the Orion-KL Hot Core or Compact Ridge. The discrepancy between the deuteration levels measured in Sgr B2(N2) and the predictions of current chemical models for complex organic molecules, and the difference between Sgr B2(N2) and Orion KL may both be due
Paper: 2019MNRAS.486.4061M (acknowledgments)
ACKNOWLEDGEMENTS This research was supported by the Netherlands Organization for Scientific Research (AAM), RFBR grant 18-502-12025 (GVL), the Royal Society (AI), the Väisälä Foundation (SST) in a part of development of a theoretical model and by the grant 14.W03.31.0021 of the Minestry of Educartion and Sciense of the Russian Federation (AAM and SST) in a part of data analyses and data comparison with the theoretical results. The authors would like to acknowledge networking support by the COST Actions CA16214 and CA16104. We are also grateful to Victor Doroshenko and Nikolai Shakura for discussion and a number of useful comments.
Paper: 2017MNRAS.466.2390C (acknowledgments)
The authors cordially thank the referee for helpful comments on the manuscript. They furthermore thank Eiichiro Komatsu for helpful discussions about the definition of f NL with respect to the commonly used WMAP parametrization. JC is supported by the Royal Society as a Royal Society URF at the University of Manchester, UK. ED acknowledges support from the DOE under grant No. de-sc0008016. MK was supported by NSF Grant No. 0244990, NASA NNX15AB18G and the Simons Foundation.
Paper: 2020MNRAS.493....1H (fulltext)
it is challenging to reproduce this effect in our simulations without additional feedback. Stars from cold accretion mostly formed at high redshifts in small haloes that later assembled into the massive galaxies. The most efficient way to remove them from our simulations is to have stronger winds, i.e. stronger mass-loading factors in those haloes. However, having too much winds early will unavoidably fail to match observations at higher redshifts. The right-hand panels of Figs 10 and 11 show that cold accretion has nearly stopped after = 2 for these galaxies but hot accretion and wind recycling continue growing rapidly at low redshifts, and is responsible for most of the excess stars formed. After = 2, the stars in massive galaxies that must be prevented from forming come from almost equal parts: hot accretion, hot wind re-accretion, and cold wind re-accretion. Hence, preventing hot gas from cooling and forming stars at these times will eliminate both the hot mode accretion and hot wind re-accretion and will lessen the tension in massive haloes. There are several potential mechanisms, such as AGN and cosmic ray heating, that could reduce the amount of this cooling gas but they are not yet included in our simulations. The hot wind recycling would be harder to affect by extra heating, because the higher metallicity of this gas makes it cool faster. In hydrodynamic simulations, mixing between ejected wind elements and the surrounding gas may have a large impact on the amounts of hot and hot wind accretion. It is, however, unclear whether this type of feedback could prevent the ∼1/3 of stars formed through cold wind re-accretion after = 2 from forming, which is also necessary to match the observations. It is possible that a more accurate treatment of the cloud-CGM interaction would allow a larger fraction of these winds to become hot, alleviating this problem. 6 THE HIGH-RESOLUTION SIMULATION OF THE REFERENCE MODEL In this section, we present key results from the high-resolution RefHres simulation. It adopts the new wind launch algorithm as described in the previous sections using our fiducial set of wind parameters listed in Table 1 . It is also implemented with the numerical improvements to the SPH hydrodynamics introduced by Huang et al. ( 2019 ). We will focus on those predictions that have changed significantly from Huang et al. ( 2019 ) and from our previously published work with our new wind algorithm. 6.1 The stellar content Fig. 14 shows that the GSMFs from our fiducial simulation, shown as black lines in each panel, are mostly consistent with observations at all redshifts from = 0 to 4. The agreement is particularly good at the faint end except for = 0, where our simulation slightly underproduces the number of these low-mass galaxies. At the massive end, our GSMFs agree with observations at = 1 and 4. However, our fiducial simulation produces too many massive galaxies at = 0 and too few massive galaxies at = 2, even after taking account
Paper: 2018ApJ...858...93I (acknowledgments)
We are grateful to Jacek Niemiec, Martin Pohl, Oleh Kobzar, Arianna Ligorini, and Artem Bohdan for fruitful discussions. Numerical computations and analyses were carried out in part on Cray XC30 and computers at Center for Computational Astrophysics, National Astronomical Observatory of Japan. This work used the computational resources of the K computer provided by the RIKEN Advanced Institute for Computational Science through the HPCI System Research Project (Project ID: hp150263), and was supported in part by JSPS KAKENHI grant No. 17H02877. This work also used the computational resources of the HPCI system provided by Information Technology Center, Nagoya University through the HPCI System Research Project (Project ID: hp170158).
Paper: 2018MNRAS.480..623C (acknowledgments)
ACKNOWLEDGEMENTS A.C., M.S. and E.Z. acknowledge support from CONICET, and E.Z. acknowledges support from SeCyT-UNC.
Paper: 2016A&A...588L...2D (fulltext)
for VLM stars and BDs have high uncertainties at ages younger than 1 Myr ( Baraffe et al. 2002 ), we assumed an age of 1 Myr for [GKH94] 41. The temperature of 2795 K then corresponds to a mass of 41 . If the uncertainty of 1.5 subclasses in the spectral type is taken into account, a possible mass range for [GKH94] 41 is from 14 to 97 . We also calculated the luminosity of the source itself (without any contribution from its surrounding envelope) to constrain the mass range. We first dereddened the 2MASS -band magnitude () using the extinction law with ( Mathis 1990 ) and the visual extinction ( Bulger et al. 2014 ). The value was obtained from the best-fit atmospheric model (see Bulger et al. 2014 and references therein) over the 2MASS bands. Using the bolometric correction in -band versus spectral type relation as given in Filippazzo et al. (2015) , we derived a luminosity of 0.025 . Figure 3 shows the luminosity versus temperature diagram for [GKH94] 41 with the theoretical evolutionary models from Chabrier et al. (2000) and Baraffe et al. (1998) . The location of [GKH94] 41 in the diagram indicates that the upper limit (including the uncertainties) to its current mass is just above the substellar limit and below 0.09 , which is consistent with the upper limit value of 97 estimated above. We therefore conclude that the upper limit to the mass of [GKH94] 41 is in the range of 1497 . Second, we estimated the envelope mass of [GKH94] 41 based on the modeling of a modified blackbody with fluxes from 70 m to mm wavelengths (Table 2 ) as used to estimate the envelope mass of IC348-SMM2E (see Palau et al. 2014 ). We obtain the dust temperature , dust emissivity and envelope mass by searching these three parameters minimizing the in the following ranges: from 10 K to 40 K, from 0.1 to 1.8 and from 1 to 60 . We assumed a dust opacity coefficient dependence on wavelength with , cm g (see Palau et al. 2014 and references therein) and a gas-to-dust ratio of 100. The best fit (Fig. 4 ) then gives = 34 K, = 2 , and = 0.4. The obtained dust emissivity index value is significantly lower than the typical value of about 1.4 for class I objects in Taurus ( Chandler et al. 1998 ). This is probably because the best fit is obtained from only three available data points. If we take = 1.4 and = 34 K, we obtain an envelope mass of 8 , a factor of 4 higher than the value estimated from the fitting. We therefore take the final mass of [GKH94] 41 to be 49 , which the sum of the current mass (41 ) and the envelope mass (8 ). If the uncertainty in the spectral type is taken into account, this upper limit will be 49. As the outflow process also ejects accreting material and dissipates material in the envelope, the final mass of [GKH94] 41 thus will very likely be below the substellar limit. Fig. 3 H-R diagram for [GKH94] 41 and IRAS 04191+1523B. Isochrones and mass tracks from the theoretical evolutionary models of Chabrier et al. (2000) for and Baraffe et al.
Paper: 2018MNRAS.477.3433E (acknowledgments)
Acknowledgements ME and HH acknowledge support from the European Research Council under FP7 grant number 279396. The authors thank Jerome Amiaux, Mark Cropper, Jean-Charles Cuillandre, Thomas Kitching, Yannick Mellier, Jason Rhodes, and Gijs Verdoes Kleijn for discussions and comments. In this project we used the scikit-learn (Pedregosa et al. 2011 ), pandas (McKinney 2010 ), scipy (Jones et al. 2001 ), matplotlib (Hunter 2007 ), and ipython (Perez Granger 2007 ) software packages. This research has made use of NASA's Astrophysics Data System.
Paper: 2019ApJ...870...20W (acknowledgments)
I appreciate support of this publication by the Solar Physics Laboratory of Goddard Space Flight Center. I thank the referee for suggesting a better ordering of the manuscript. I am grateful to PMOD/WRC, Davos, Switzerland, for supplying their irradiance record (Frohlich 2006 ), composite_42_65_1709.dat, with new data from Virgo on the ESA/NASA Satellite Mission, SOHO .
Paper: 2020A&A...633A.144V (acknowledgments)
The authors would like to thank J. D. P. Kenney, E. J. Murphy, and Olaf Wucknitz for useful comments on the article.
Paper: 2017MNRAS.471.1330A (fulltext)
fluencies (chemical equilibrium). The chemical equilibrium during ice processing by photons, electrons or ions occurs at large fluences when parent and daughter species reach abundances that keep each other's abundance constant during reciprocally production-destruction reactions in the presence of incoming energy. This steady state scenario explains the horizontal plateau in chemical abundances observed at fluences larger than ∼1.4 × 10 13 ion cm −2 (see for example Fig. 4 ). Additionally, in the case of ion bombardment, the sputtering may affect drastically the molecular abundances as function of time, mainly in the beginning when parent species are more affected by sputtering. At high fluences, the molecular abundance reaches the chemical equilibrium, the sputtering affects equally all the species and their column densities decrease (in a constant rate) as function of fluence. The chemical equilibrium is also affected by the sample temperature since volatile products tend to desorb from the ice. More details on chemical equilibrium induced by radiation in ices are given elsewhere (Vasconcelos et al. 2017 ). Studying the radiolysis of N 2 : CO (1:1) ice at 10 K by electron beam irradiation with final energy dose similar to the exposure of astrophysical ice to 5 × 10 6 year in the ISM, Jamieson et al. ( 2005 ) have found the result N 2 O/N 2 > 6.0 × 10 −4 . Bombarding N-rich ices by protons and considering the band strengths in the infrared for N 2 O and N 2 derived by Jamieson et al. ( 2005 ), Loeffler et al. ( 2010 ) have found N 2 O/N 2 ∼ 2.0 × 10 −3 for the equilibrium chemistry attainted at longer fluencies. This value is very similar to the one observed in the gas phase at Sgr B2, as shown above (Womack, Ziurys Wyckoff 1992 ). The current manuscript investigates the survival of solid N 2 O against ionizing field (1.5 MeV N + cosmic ray analogues). From the measured destruction cross-section (σ d ) and the ion flux (Φ) in the laboratory, the photodissociation rate at laboratory (k d = σ d × Φ), turns out to be about 7.0 × 10 −4 ions s −1 . Same equation can be employed to derive the photodissociation rate in space. Considering first order kinetics for the destruction of N 2 O by swift ions, the evolution of N 2 O abundance as a function of time is described by equation ( 25 ) where N(t) and N(t 0 ) represents the molecular abundance at a given time (t) and at the beginning of experiment. κ d is the dissociation rate in units of ions s −1 , σ d is the dissociation cross-section in units of cm 2 and Φ is the ion flux in units of ions cm −2 s −1 . To extrapolate the current laboratory results to astrophysical environments, the flux of medium-mass ions (3 Z 11) having energies between 0.1–10.0 MeV/u estimated by Pilling et al. ( 2013 ) is used. The values considered are 3.0 × 10 −1 and 1.0 × 10 −1 ion cm −2 s −1 , for the ISM and inside the Solar system at Pluto orbit, respectively. Fig. 5 presents the evolution of N 2 O abundance, for ices
Paper: 2016ApJ...833....5S (acknowledgments)
This work is dedicated to the memory of our friend and colleague, Mike Bradford. SDO data are provided courtesy of NASA/ SDO and AIA and HMI science teams. EUVI data are provided courtesy of the STEREO Sun Earth Connection Coronal and Heliospheric Investigation (SECCHI) team. FIRS has been developed by the Institute for Astronomy at the University of Hawai‘i, jointly with the National Solar Observatory (NSO). The FIRS project was funded by the National Science Foundation Major Research Instrument program, grant number ATM-0421582. T.A.S. gratefully acknowledges colleagues at the High Altitude Obseratory for interesting discussions regarding this work during a scientific visit. This work also benefited from discussions at the International Space Science Institute (ISSI) meetings on “Sub-arcsec Observations and Interpretation of the Chromosphere” and “Implications for Coronal Heating and Magnetic Fields from Coronal Rain Observations and Modelling.” This research has made use of NASA’s Astrophysics Data System. The authors thank the referee for a careful review and fruitful comments. Facilities: Dunn(FIRS) , SDO (AIA) , SDO (HMI) , STEREO (EUVI-A). Software: SSWIDL, SolarSoft, Crispex (Vissers Rouppe van der Voort 2012 ), HeLIx + (Lagg et al. 2004 ; Lagg 2007 ), PIKAIA (Charbonneau 1995 ).
Paper: 2018ApJ...861...50B (fulltext)
between and M * in the top panel of Figure 1 for isolated galaxies with ALFALFA data and the residuals () between the observed and the predicted in the bottom panel. Equation ( 2 ) is plotted as the solid curve. The dashed lines in the bottom panel are ±1 dex from the relation, which roughly corresponds to the 3 σ scatter. We use the −1 dex residual as a reference point throughout the paper as an arbitrary line of demarcation. Figure 1. Top panel: the isolated to M * relation of gas-rich galaxies. Only isolated galaxies with detections from the 70% ALFALFA survey are plotted here as black contours and dots. The fit from Equation ( 2 ) is plotted as the solid curve and noted in the legend. Bottom panel: H i mass residual ( −1 dex) as a function of stellar mass with the same symbols at the top panel. The ±1 dex (3 σ ) residuals are shown as dashed lines. 2.4. Emission Line Measurements Optical emission line flux ratio diagnostics are often used to differentiate between various physical phenomena that ionize the ISM of galaxies Baldwin et al. ( 1981 ), Veilleux Osterbrock ( 1987 ). Differentiating between excitation sources (actively accreting black holes, star formation, shocks from AGB stars, etc.) depends on the ionization mechanism strength. Reines Volonteri ( 2015 ) measure emission line strengths of ∼67000 galaxies selected from the NSA catalog. The authors cull the spectral catalog with an S/N cutoff of three and an equivalent width cutoff of one for the H α , [N ii ], [O iii ] emission lines and an S/N cutoff of two for the H β line. They iteratively subtract continuum and absorption lines using the models of Tremonti et al. ( 2004 ) from the observed spectrum, carefully masking out the emission lines. Next, a one- or two-component Gaussian is fit to the [S ii ] doublet and used as a template for the remaining emission lines of interest, excluding [O iii ]. We use the data set from this study without any modifications. 2.5. BPT Emission Line Diagnostics We examine the emission line diagnostics of all galaxies from Reines Volonteri ( 2015 ), regardless of environment or H i data using the [O iii ]/H β and [N ii ]/H α narrow-line BPT diagram (Baldwin et al. 1981 ; Veilleux Osterbrock 1987 ; Kewley et al. 2006 ). We adopt the division into the star-forming, composite, and AGN categorization that has been performed by Reines Volonteri ( 2015 ) using the Kewley et al. ( 2001 ) and Kauffmann et al. ( 2003 ) categorization. The Kewley et al. ( 2001 ) demarcation line represents the division between a theoretical maximal starburst model and emission that can only be explained by AGN activity. The Kauffmann et al. ( 2003 ) demarcation line is an empirical separation between the star-forming sequence and the AGN sequence. Galaxies below the Kauffmann et al. ( 2003 ) line are considered to be purely star-forming and galaxies above the Kewley et al. ( 2001 ) line are considered Seyferts and LINERs. The region between these two demarcation lines
Paper: 2019ApJ...877...63S (acknowledgments)
Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS website is http://www.sdss.org/. This publication makes use of data products from the Wide-field Infrared Survey Explorer , which is a joint project of the University of California, Los Angeles, and the Jet Propulsion Laboratory/California Institute of Technology, and NEOWISE , which is a project of the Jet Propulsion Laboratory/California Institute of Technology. WISE and NEOWISE are funded by the National Aeronautics and Space Administration.
Paper: 2015MNRAS.449..933A (acknowledgments)
This work is supported by the Scientific and Technological Research Council of Turkey (TÜBİTAK) under the grant 113F354. MAA is a member of the Science Academy (Bilim Akademisi), Turkey. We thank Patrick Weltevrede for kindly providing us with unpublished post-glitch timing data, and Danai Antonopoulou, Anna Watts, Cristobal Espinoza and Patrick Weltevrede for useful discussions and for sharing with us their preprint with an alternative approach to the post-glitch behaviour of PSR J1119−6127. We thank the referee for useful comments.