adsabs/WIESP2022-NER huggingface dataset:
Paper: 2019ApJ...875...36S (fulltext)
of the cavity λ H decreased due to subtraction of solar wind He + ions from ∼19.5 to ∼15.5 au, whereas the average ionization factor β H,E decreased from 7.6 × 10 −6 to 3.9 × 10 −6 s −1 . These values are still well above the expected range of 3–6 au for λ H and (0.5–1) × 10 −6 s −1 for β H,E . Therefore, the presence of solar wind He + ions in SWAP observations does not fully explain this discrepancy, but it is an important factor that influences the fitted spectra of H + PUIs. Moreover, a reasonable agreement may not be possible using the Vasyliunas Siscoe ( 1976 ) model, which is only a simple approximation and does not account for most of the transport processes. More advanced models (see, e.g., Zank et al. 2018 ) can give results consistent with the observations for more realistic parameters. 6. Summary Helium, the second most abundant species in the solar wind, is emitted from the solar corona almost entirely in the form of double ionized alpha particles. Due to charge exchange with ISN atoms, some of the He 2+ ions acquire electrons and form a new population of singly charged solar wind He + ions. Charge exchange with ISN hydrogen is the dominant source of this population. The rate of this reaction significantly increases over the range of the solar wind speeds. As a result, the ratio of solar wind He + ions to solar wind He 2+ can exceed 10% at the termination shock for regions of fast solar wind ≥600 km s −1 , such as at high heliolatitudes around solar minimum (McComas et al. 2000 ). Therefore, solar wind He+ ions are an important component of the plasma and, as such, should be included in analyses of the processing of the solar wind and acceleration processes at the termination shock. In this study, we show that this contribution is visible in some part of the SWAP observations as a “bump” just below the cutoff energy of H + PUIs. This paper is a first report on the observation of this population in the outer heliosphere. Earlier reports on solar wind He + ions are very sparse (Tucker Gould 1966 ; Kozlovsky 1968 ; Feldman et al. 1974 ; Gloeckler Geiss 1998 ). We select observations with a local maximum just below the H + PUI cutoff energy. About ∼3/5 of them are related to significant changes in the solar wind speed up to 10 days before the date of the observation. Consequently, the observed spectrum was likely modified due to this change. In the remaining ∼2/5 cases, the solar wind speed remains stable throughout 10 days, and for most of them, we find significant improvements in the goodness of the fit for H + PUIs when solar wind He + ions are adequately accounted for. In only one observation is the improvement not statistically significant. We point out that the observations of solar wind He + ions are more likely to be visible during times of faster solar wind, with lower temperatures and high relative abundances of solar wind alpha particles. We notice that in some observations also, other heavy ions such as oxygen and carbon
Paper: 2020A&A...641A..60P (fulltext)
rot / σ profiles tend to peak at a median radius of R ∼ 3 R e . By comparison, the ePN.S FRs tend to peak at smaller fractions of R e , at a median 1.3 R e . This difference between the observed and simulated FRs is not a consequence of the selection functions of the samples. The TNG galaxies are selected according to color, mass and intermediate to major axis ratio p . The selection in p removes centrally elongated galaxies, most of which have intermediate to low λ (1 R e ) (Fig. 5 ). The Atlas3D sample, selected as described in Sect. 3 , has some of the disk galaxies removed which as in MANGA (Fig. 5 ) will be mostly located at high λ (1 R e ). We recall that in the comparison we consistently matched the color selection and mass range of the Atlas3D sample to the selection criteria adopted for the TNG galaxies. Thus the TNG sample should in principle contain a larger number of late type galaxies with strong disks, i.e. high V rot / σ , by comparison with Atlas3D. Figure 13 shows instead that the TNG100 and TNG50 ETG samples lack galaxies with high rotational support at 1 R e . In Sect. 6.1 we discussed that the effective radii used to normalize the radial scales in TNG would be expected to be systematically slightly overestimated compared to the observed R e since they are defined using the total bound stellar mass which is often inaccessible in observations. If taken into account, this effect would increase the gap between simulated and observed samples. On the other hand, since the mass-size relation is roughly reproduced in the simulations (Fig. 7 ), the different steepness of the V rot / σ profiles in observations and simulations implies a different distribution of the angular momentum as a function of radius in the simulated galaxies. This could be due to a too efficient condensation of the gas into stars that does not allow the gas to dissipate and collapse to small enough radii. 7.3. Simulated versus observed profiles – outskirts Figure 14 compares the relation between V rot / σ in the central regions and V rot / σ in the stellar halos for the simulated and observed galaxies. The observed galaxies are from the ePN.S survey ( Pulsoni et al. 2018 , their Fig. 9) and are reported in the top left panel. Their measurements use absorption line data at 1 R e and PN data for the halos, which on average cover 6 R e with a large scatter (minimum 3 R e , maximum 13 R e ). The central and right top panels show V rot / σ (6 R e ) versus V rot / σ (1 R e ) for TNG100 and TNG50 ETGs, respectively. Galaxies close to the dashed 1:1 lines show similar rotational support in the central regions and in the outskirts; galaxies below the lines have increased rotational support in their stellar halos; galaxies above the lines instead have reduced rotation at large radii. Fig. 14. V / σ ratio in the central regions compared with the V / σ at large radii. Top : V / σ (1 R e ) versus V / σ (6 R e ) for the ePN.S galaxies ( left ), TNG100 ( middle ), and TNG50 (
Paper: 2016A&A...585A.150N (fulltext)
average confirm the continuous decline of metallicity. The use of age groups in the inner disc revealed an increase of the metallicity level with age. The older clusters (1.7 Gyr) appear more metal rich by about 0.07 dex than clusters of a young group (0.5 Gyr). This is clearly the opposite of what was commonly looked for in the past, a decline of metallicity with age, but it might be an observational confirmation of radial migration. However, to further validate this result, a homogeneous age scale for the clusters and a somewhat larger sample of older objects are needed. Furthermore, we compared open clusters and Cepheids with recent Galactic chemical and chemodynamical models. The models reproduce the observed increase of metallicity with age and generally agree well with the observations of young and older populations. However, the comparison shows that the models probably underestimate the effect of radial migration, or that there is another not yet included mechanism acting in the outer disc, where open clusters show a much flatter gradient than predicted by the models. However, to study the reliability of the current models in more detail and to further improve them, two questions apparently need to be resolved first: a somewhat larger open cluster sample that covers a broader Galactic distance range in each age group is needed. Most studied objects are still concentrated in the solar circle. Second, a more detailed investigation of the relative metallicity scale between open clusters and Cepheids is required. This should include an improvement of the knowledge of the metallicities of Cepheids that are cluster members and of the metallicities of the hosting open clusters, but the latter based on photometric methods or spectroscopic data of red giants or main-sequence stars. Finally, we showed the Galactic metallicity distribution of a 8 kpc large area in a contour plot and the metallicity as a function of the azimuthal angle. Both confirm that there is little difference between the Galactic quadrants. However, there are some slightly overabundant areas that might be related to the spiral structure of the Galaxy. The final results of spectroscopic surveys, such as the Gaia -ESO Public Spectroscopic Survey or the Gaia mission itself, will certainly provide the basis for more detailed studies. Furthermore, the availability of current and future homogeneous photometric datasets will allow complementing the spectroscopic metallicities with results obtained by photometric methods. We wish to mention in particular the DG method, which can be applied to any photometric data with available temperature calibrations. Appendix A: Additional table Table A.1 Metallicities of the open clusters. Notes. The HQS, LQS, and weighted photometric metallicities, respective errors, number of stars used for the calculation (N), and the total weight (W) of the photometric metallicities. The code for the photometric systems is 1 (DG), 2 (DDO), 3 (Friel), 4 (Ca
Paper: 2015ApJ...804...55L (acknowledgments)
We thank the anonymous referee for helpful comments that greatly improved the presentation of this paper. We also thank H.J. Mo and Frank C. van den Bosch for useful discussions. This work is supported by the 973 Program (No. 2015CB857002), NSFC (Nos. 11128306, 11121062, 11233005), the Strategic Priority Research Program “The Emergence of Cosmological Structures” of the Chinese Academy of Sciences, Grant No. XDB09000000, and a key laboratory grant from the Office of Science and Technology, Shanghai Municipal Government (No. 11DZ2260700).
Paper: 2019ApJ...875...96N (acknowledgments)
We are grateful to Professor Inga Kamp and Dr. Satoshi Okuzumi for their useful comments. We thank the referee for many important suggestions and comments. This paper makes use of the following ALMA data: ADS/JAO.ALMA #2015.1.01259.S and ADS/JAO.ALMA #2013.1.00601.S. ALMA is a partnership of European Southern Observatory (ESO) (representing its member states), National Science Foundation (USA), and National Institutes of Natural Sciences (Japan), together with National Research Council (Canada), National Science Council and Academia Sinica Institute of Astronomy and Astrophysics (Taiwan), and Korea Astronomy and Space Science Institute (Korea), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, Associated Universities, Inc/National Radio Astronomy Observatory (NRAO), and National Astronomical Observatory of Japan. Our numerical studies were carried out on SR16000 at Yukawa Institute for Theoretical Physics (YITP) and computer systems at Kwasan and Hida Observatory (KIPS) in Kyoto University, and PC cluster at Center for Computational Astrophysics, National Astronomical Observatory of Japan. ALMA Data analysis was carried out on the Multi-wavelength Data Analysis System operated by the Astronomy Data Center (ADC), National Astronomical Observatory of Japan. S.N. was supported by the ALMA Japan Research Grant of NAOJ Chile Observatory, NAOJ-ALMA-211. This work is supported by JSPS (Japan Society for the Promotion of Science) Grants-in-Aid for Scientific Research (KAKENHI) (grants No. 25108004, 25108005, 25400229, 15H03646, 15K17750, 17K05399, and 17K14244), by Grants-in-Aid for JSPS fellows (grant No. 16J06887), and by the Astrobiology Center Program of National Institutes of Natural Sciences (NINS; grant No. AB281013). C.W. acknowledges support from the Science and Technology Facilities Council (STFC; grant No. ST/R000549/1) and start-up funds from the University of Leeds. Astrophysics at Queen’s University Belfast is supported by a grant from the STFC (ST/P000312/1). T.J.M. thanks Leiden Observatory for hospitality.
Paper: 2018MNRAS.481.3108C (acknowledgments)
ACKNOWLEDGEMENTS We thank Volker Springel for access to arepo . We thank Joel Primack for insightful questions about the spin parameter that motivated this project. We also thank Ramesh Narayan, Re’em Sari, Rachel Somerville, James Bullock, and Michael Boylan-Kolchin for enlightening discussions. We also thank Cristina Popa for her help in the early stages of this project.
Paper: 2016A&A...587A..58B (fulltext)
2011a ; Konopacky et al. 2013 ; Lee et al. 2013 ; Barman et al. 2015 ). The good fit in the 35 m region for the planets d and e (along with the other wavelengths), despite the lack of non-equilibrium chemistry, non-solar metallicity compositions, and local variation of the cloud thickness in the models, suggest that these ingredients are not critically needed here. Our analysis rather demonstrates that a higher content in dust and reduced CIA of H resulting from a low surface gravity can explain the deviation of the planet’s SED with respect to those of standard field L and T dwarfs. This does however not mean that such additional ingredients do not play a role, especially for HR8799b and c. Barman et al. (2011a) noted that the -band spectrum of the high-amplitude variable brown dwarf T1 dwarf 2MASS J21392676+0220226 fitted well the SED of HR8799 b. The spectrophotometric properties of the two objects representing the best the properties of HR8799 b and c, 2MASS J13243553+6358281 and SDSS J151643.01+305344.4 , can be modeled by blended spectra of L-T dwarfs, e.g. as unresolved L/T transition binary at field dwarf ages ( Burgasser et al. 2010 ; Geißler et al. 2011 ). Beside its possible young age, we can speculate that SDSS J151643.01+305344.4 is a good template because the integrated flux of the possible binaries mimics layered patchy atmospheres. Conversely, Heinze et al. (2015) find that 2MASS J13243553+6358281 experiences high amplitude variability in the red-optical. The planets e and d could also be variable objets. The best fitting template to HR8799d’s SED – PSO J318.5338-22.8603 – was recently found to experience photometric variability in the near-infrared ( Biller et al. 2015 ). In addition, the red late-L dwarf WISE 1738, which reproduces well the properties of HR8799d and e also has a highly variable near-infrared spectrum ( Mace et al. 2013 ). In that sense, the good fit of the planet spectra with field dwarfs templates reddened by sub-micron iron, corundum, or silicate grain opacities in Sect. 3 is in line with the conclusions of Yang et al. (2015) who find that the observed variability of L5 dwarfs can be explained by the presence of spatially varying high-altitude haze layers above the condensate clouds. Such a layer may exist into the planet atmospheres and could be added to the current self-consistent cloud models (e.g. BT-SETTL, DRIFT-PHOENIX) to bring the and radii estimate in agreement with the expectations from evolutionary models. 5.2. Masses and diversity among the HR8799 system Marois et al. (2008) , Marois et al. (2010) , and Currie et al. (2011b) used the bolometric luminosity of the planets to estimate masses of 5, 7, 7, and 7 for HR8799b,c,d and e, respectively if the system is 30 Myr old. We retrieve the same luminosity – and then mass estimate – for HR8799d and e using the bolometric correction estimated for the dusty L7 dwarf WISE 0047 by Gizis et al. (2015) . The radii and log derived from the best atmospheric
Paper: 2019ApJ...887..208P (acknowledgments)
O.P. sincerely thanks Dr. A. Retinò and Prof. W. H. Matthaeus for the fruitful discussions. O.P. and S.S. are partly supported by the International Space Science Institute (ISSI) in the framework of the International Team 405 entitled “Current Sheets, Turbulence, Structures and Particle Acceleration in the Heliosphere.” D.P. was partially supported by STFC grant ST/N000692/1. Numerical simulations discussed here have been performed on the Marconi cluster at CINECA (Italy), within the projects IsC53_RoC-SWT and IsC63_RoC-SWTB, and on the Newton cluster at the University of Calabria (Italy). This project (F.V., S.S.) has received funding from the European Unions Horizon 2020 research and innovation program under grant agreement No. 776262 (AIDA, www.aida-space.eu ).
Paper: 2018A&A...616A..48S (fulltext)
to the general population of XRT light curves of GRBs. The gray shaded region is a compilation of long GRB light curves ( Evans et al. 2007 , 2009 ), where the color represents density. The light blue lines are sGRB light curves from bursts with duration of T 90 ≲ 2 s and those that were classified as short in Kann et al. (2011) , Berger (2014) and D’Avanzo et al. (2014) . The thick green line is GRB 111117A. Despite the remarkably high redshift, the luminosity is comparable to the bulk of the short burst population, and subluminous compared to the lGRB population. The inset shows the X-ray luminosity distributions of sGRBs and lGRBs at 10 ks, indicated by the vertical dashed line in the main panel. The separation of GRBs in two distinct classes based on their high-energy observables points to an intrinsically different physical origin. Long GRBs are typically interpreted as collapsars ( MacFadyen Woosley 1999 ) in which a single, massive star undergoes gravitational collapse, where the currently preferred model for sGRBs is the merger of two NSs ( Eichler et al. 1989 ; Nakar 2007 ). Bromberg et al. (2013) investigated the degree to which high-energy observables of the long and short GRB populations overlap, and quantified the certainty in class membership. According to Bromberg et al. (2013) , GRB 111117A has a probability of being a sGRB. Compared to two other sGRB candidates at high redshift, GRB 060121 ( T 90 = 1.97 ± 0.06 s; de Ugarte Postigo et al. 2006 ; Levan et al. 2006 ) at 1.7 ≲ z ≲ 4.5 () and GRB 090426 ( T 90 = 1.28 ± 0.09 s; Antonelli et al. 2009 ; Levesque et al. 2010 ; Thöne et al. 2011 ) at z = 2.609 (), the certainty in class membership for GRB 111117A is much higher. Additionally, Horváth et al. (2010) classify both GRB 060121 and GRB 090426 as intermediate-duration bursts because both events have very soft spectra, compared to the hard ones typically seen in sGRBs. Intermediate bursts are very clearly related in their properties to lGRBs ( de Ugarte Postigo et al. 2011 ), so they are unlikely to come from compact object mergers. GRB 111117A is also securely classified as a sGRB according to the Horváth et al. (2010) classification scheme. A number of other GRBs are thought to be short and at relatively high redshift ( z ≳ 0.9). If we consider bursts with a probability of being short f NC > 50%, according to the Bromberg et al. (2013) classification scheme, five sGRB are found: GRB 051210 at z ~ 1.3 ( f NC = ; Leibler Berger 2010 ), GRB 060801 at z = 1.131 ( f NC = ; Berger et al. 2007 ), GRB 070714 at z = 0.923 (no f NC due to extended emission; Graham et al. 2009 ), GRB 090510 at z = 0.903 ( f NC = ; McBreen et al. 2010 ), and GRB 100117 at z = 0.915 ( f NC = ; Fong et al. 2011 ). Although in individual cases a secure host association (hence redshift determination) is uncertain, there does seem to be a number of sGRBs at z ~ 1. This certainly makes GRB 111117A, by far, the highest redshift sGRB detected to date. The
Paper: 2015ApJ...813...54T (acknowledgments)
We thank Nick Battaglia, Rychard Bouwens, Camila Correa, Jeremy Tinker, and Rogier Windhorst for helpful discussions. We thank Rick Costa, Roberto Gomez, and Bill Wichser for help with computing. H.T. acknowledges support from NASA grant ATP-NNX14AB57G and NSF grant AST-1312991. R.C. acknowledges support from NASA grant NNX12AF91G. Simulations were run at the NASA Advanced Supercomputing (NAS) Division, Pittsburgh Supercomputing Center (PSC), and the Princeton Institute for Computational Science and Engineering (PICSciE).
Paper: 2016A&A...592A.124G (acknowledgments)
This work was supported by “Programme National de Physique Stellaire” (PNPS) and “Programme National de Physique Chimie du Milieu Interstellaire” (PCMI) from INSU/CNRS. WV’s research is funded by an ERC starting grant (3DICE, grant agreement 336474). D.S. acknowledges support by the Deutsche Forschungsgemeinschaft through SPP 1385: “The first ten million years of the solar system a planetary materials approach” (SE 1962/1-3)” This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France.
Paper: 2018A&A...615L..13G (fulltext)
as Giles et al. (2018) . A pipeline was built for the CORALIE spectra based on iSpec 1 ( Blanco-Cuaresma et al. 2014a ). All observations were aligned and co-added to increase the signal-to-noise ratio (S/N), were reduced and spectrally fitted using the code SPECTRUM ( Gray Corbally 1994 ) as the radiative transfer code. Atomic data were obtained from the Gaia -ESO Survey line list ( Heiter et al. 2015b ). We selected the line based on an R ~ 47 000 solar spectrum ( Blanco-Cuaresma et al. 2016 , 2017 ), and we used MARCS model atmospheres ( Gustafsson et al. 2008 ). The resulting errors were increased by quadratically adding the dispersions found when analysing the Gaia benchmark stars ( Heiter et al. 2015a ; Jofré et al. 2014 ; Blanco-Cuaresma et al. 2014b ) with the same pipeline. This resulted in an effective temperature of 4877 ± 68 K, a log g of 3.41 ± 0.07 dex, and [Fe/H] = −0.24 ± 0.04 dex. In the second data release of Gaia ( Gaia Collaboration 2018 ), EPIC248847494 has a measured parallax (see Table 2 ) based on which we can determine an independent stellar radius using bolometric absolute magnitudes and the spectroscopically determined effective temperature for EPIC248847494 following the method detailed in Fulton Petigura (2018) . We took the K -band apparent magnitude ( Skrutskie et al. 2006 ), the Gaia distance, and a bolometric correction (BC K , from Houdashelt et al. 2000 ) of 1.91 ± 0.05, which was interpolated from the range within the coarse grid. We chose not to include an extinction correction as this only introduces an uncertainty of 0.5% ( Fulton Petigura 2018 ). This gave a radius of 2.70 ± 0.12 R sol . Table 2. Properties of the EPIC248847494 system. Notes. (a) Huber et al. (2016) , (b) APASS: Henden Munari (2014) , (c) 2MASS: Skrutskie et al. (2006) , (d) Gaia Collaboration (2018) , (†) This Work. Taking the spectrally determined metallicity and effective temperature and the measured radius as observational constraints, we input them into the Geneva stellar evolution code ( Eggenberger et al. 2008 ). This resulted in a stellar mass of 0.9 ± 0.09 M sol . These values of mass and radius would therefore indicate a log g of 3.52 dex. When we fixed the iSpec analysis to this log g , the metallicity and effective temperature were very similar to the initial results (see Table 2 ). Log g is not well constrained spectroscopically, and changes have a very limited effect on other parameters. Therefore we adopt the parameters based on log g = 3.52. 3.2. Eliminating the photometric systematics of K2 The possibility for false positives is high in monotransits. We therefore endeavored to eliminate all causes for false positives. All objects listed as “stars” with K2 light curves within 25 arcmin were checked for similar artefacts. Of the 61 objects, none showed odd behaviour at the same epoch as the monotransit. Additionally, the location of EPIC248847494 was not near the edge of the CCD, which suggests that no near-edge effects
Paper: 2015MNRAS.452..333C (acknowledgments)
This work was carried out using the computational facilities of the Advanced Computing Research Centre, University of Bristol. 1 PC thanks the University of Bristol for a Postgraduate Research Scholarship. We thank the referee, Erin Kara, for very useful comments which have improved the paper.
Paper: 2019A&A...625A..37N (acknowledgments)
The authors are thankful to A. A. Osipova (Pulkovo Observatory) for preparing sunspot data for Table 1 . AN and LK acknowledge the support by the Russian Foundation for Basic Research (project 17-52-80064_BRICS) and by budgetary funding of Basic Research programme II.16. DB acknowledges the support from DST, India for BRICS project grant.