adsabs/WIESP2022-NER huggingface dataset:
Paper: 2019ApJ...872...74K (fulltext)
plotted the intensity increase and decrease in Figure 6 (a). The base-subtracted images show that strong coronal dimming appears in the west of the active region from 03:10 UT, which is 18 minutes after the GOES peak time. The dimming reached the peak at 04:10 UT, which is 78 minutes after the GOES peak time. Coronal mass ejections (CMEs) are often associated with coronal dimming (Sterling Hudson 1997 ). Therefore, in this case, a CME was likely to occur as well, even though the CME catalog 3 3 https://cdaw.gsfc.nasa.gov/CME_list/ by the SOHO /Large Angle and Spectrometric COronagraph (LASCO; Brueckner et al. 1995 ) does not show any CME. Figure 6. (a) Spatially integrated intensity increase (thick black) and decrease (dashed red) of the 195 Å intensity measured on EIS 195 Å images from 01:46:55 UT. The vertical dotted line indicates the GOES peak time at 02:52 UT. (b) Intensity and base-subtracted maps of EIT 195 Å. The observation time is shown at the top right of each image. The base-subtracted images are the difference from the image at 01:46:55 UT shown in the top plot. The red boxes show the field of view of Figure 3 . 5. Summary and Discussions In this section, we summarize our results from the statistical and event studies and discuss the chance to observe the linear polarization in H α and its origin. 5.1. Infrequent Occurrence of Significant Linear Polarization at H<italic>α</italic> in Solar Flares We found only one event that showed significant linear polarization signals among 71 H α flares including 64 GOES flares with the spatial and temporal samplings of 10′′ × 10′′ and 40 s, respectively. Many preceding papers reported the characteristics of the linear polarization in H α only for the events that show significant polarization signals. In the cases where the polarization was not detected, it has been considered that the sensitivity of the polarimetry was insufficient. On the other hand, Bianda et al. ( 2005 ) reported that they did not find significant linear polarization signals in 30 flares, even though they used a high-sensitivity instrument. Based on their preliminary analysis of dozens of flares, Hanaoka ( 2004 , 2005 ) also found that one flare showed significant linear polarization signals, while most other flares showed none. In our study, we used the polarimeter with higher accuracy described in the papers by Hanaoka ( 2004 , 2005 ) and confirmed their preliminary results. Furthermore, although the analyzed samples include some large flares such as 4 X- and 13 M-class events, the results are basically negative. The high-frequency modulation techniques reduce spurious signals caused by the seeing effect and by flaring structures evolving in time, and data obtained by these techniques are more reliable than data that were taken by low-frequency modulation techniques. Thus, our results proved that the linear polarization does not appear in most of the flares, but there are a few flares that actually show the polarization.
Paper: 2016ApJ...828...49H (fulltext)
length of just 0.5 kpc, the gas will have cooled and recombined before escaping. These baryons will therefore become visible to Lyman series absorption studies and the oxygen will be visible to O i absorption (e.g., at 1302 Å). APPENDIX MATCHING THE POINT-SPREAD FUNCTION The diffuse halo of O vi emission could be mimicked if the F125LP PSF is broader than that of F140LP. We used custom software to match images to a common PSF. Point sources (stars) that are bright in the FUV are rare, making empirical studies of the SBC PSFs challenging, and the models available from STScI severely underestimate the power in the extended wings of the PSF (at radii larger than 1 arcsec). We instead build empirical PSF models by examining multi-epoch calibration data of the globular cluster NGC 6681. These were obtained under many programs by the ACS instrument team at STScI, stretching back to 2002 when the instrument was first mounted. We have obtained all the calibration images taken in F125LP and F140LP from the MAST, drizzled and stacked them using the same methods as for our science data. The total exposure time in F125LP is 13,530 and 7350 s in F140LP. The brightest star in NGC 6681 has a FUV AB magnitude of 18.0, giving a countrate of 22 counts per second in the brightest pixel in F125LP. This is just below the nonlinear regime of the SBC detector, indicating that calibration data are as good as possible for deriving the PSF. We then selected 45 (48) stars in the F125LP (F140LP) images, extracted stamps, and stacked them to produce a very deep point-source image. A model PSF consisting of a combination of three 2D functions (an elliptical Moffat function, a radially symmetric Gaussian function, and a symmetric fourth-order polynomial function) was fitted to the stacked image using least-squares fitting with the LMFIT package for Python. 6 6 Newville et al., LMFIT: Non-Linear Least-Square Minimization and Curve-Fitting for Python, 2014, doi: 10.5281/zenodo.11813 . The stacking was then redone iteratively with the contribution from nearby stars removed by subtracting the model PSF. The final stacked images in the two filters were normalized to have a total sum of 1 and are presented in Figure 12 , while the final best fitting models were used to compute the matching convolution kernel. Figure 12. Point-spread function of the SBC filters. The top left panels show the empirical PSF of the F125LP (left-hand side) and F140LP (right-hand side) filters in log scale. The bottom left panels show the relative difference map, (F125LP–F140LP)/F140LP, without any PSF matching on the left and with PSF matching on the right. Note that the 2D PSF models are only trustworthy (at a level) out to a radius of about 1.3 arcsec. The top right panel shows the surface brightness profiles of the two empirical PSFs with the shaded regions showing the errors per annulus, and the inset showing the inner 0.5 arcsec of the profile. The lower right panel shows the signal-to-noise
Paper: 2019MNRAS.483.4766P (fulltext)
dex) that led us to speculate about the possibility that the assumption for young MC clusters to have metal contents similar to the mean galaxy present-day metallicities is justified for statistical purposes. Otherwise, when young clusters are studied to search for chemical abundance anomalies, light element abundance variations, binary fraction, extended main sequence turnoffs, among others, the knowledge of their actual metallicities could have an impact. This could be the case, for instance, of NGC 1844, NGC 1847, and NGC 330, for which Milone et al. ( 2013 ), Niederhofer et al. ( 2015 ) and Milone et al. ( 2018 ) adopted, respectively, more metal-rich metal abundances to show evidence of multiple populations. In order to see whether a better tracking of the split main sequences can be achieved, it would be worth trying to match their CMDs with theoretical isochrones with metallicities similar to those derived in this work. In the case of NGC 1850, a cluster with a large population of near-critically rotating stars, a slightly more metal-rich value has usually been adopted (Bastian et al. 2017 ). The existence of a spread in metallicity within the younger stellar populations of both MCs is well known. Piatti and Geisler ( 2013 ), using an homogeneous age/metallicity compilation, showed that the FWHM of such a scatter is 0.51 dex for the LMC and 0.32 dex for the SMC (see also Choudhury, Subramaniam Cole 2016 ; Choudhury et al. 2018 ); the MC cluster populations also exhibit a noticeable scatter at their younger end (see also Perren, Piatti Vázquez 2017 ). In this context, most of the present studied clusters are within the expected metallicity range, while others fall at the edge of the known metallicity distributions. This is the case of NGC 330, the most metal-poor young SMC cluster ([Fe/H] = −1.15 dex) so far. In the LMC, NGC 1847 resulted to be the most-metal poor young cluster ([Fe/H] = −0.91 dex) ever known, while NGC 1711 turned out to be at the metal-rich end of the LMC cluster’s metallicities ([Fe/H] = −0.06 dex). These relatively extreme metallicity values tell us that the gas out of which these young clusters were formed was not well mixed. Note that there has been speculations about the role of infall of unenriched (or less enriched) gas into the MCs leading to an unexpectedly large spread in cluster abundances at a relatively constant age (see e.g. Da Costa Hatzidimitriou 1998 ). We looked at the cluster positions in their respective host galaxies in order to search for any link of their derived metallicities with the chemical evolution histories of the MCs, particularly of those with more extreme values. Fig. 2 depicts with black points the spatial distributions in both MCs of all the clusters catalogued by Bica et al. ( 2008 ). The studied clusters are drawn with big filled circles. As a spatial reference, we have also included the areas defined by Harris and Zaritsky ( 2009 ) in the LMC main body (the bar is traced with a
Paper: 2021A&A...648L...2P (fulltext)
in denser environments, for example, star clusters ( Varri et al. 2018 ) or the Galactic bulge, where thousands of NSs may be present (e.g. Fragione et al. 2018 ). The encounter velocities originating from SN kicks will be significantly higher than the virialised velocities that have been used in published simulations (e.g. Wang et al. 2020 ). In fact, they are closer to what would be expected for a high-velocity star in the Solar neighbourhood ( Laughlin Adams 2000 ). An estimate of the gravitational capture distance, assuming that the intruder can be bound in the centre of mass frame, is given by . This scales as where ‘I’ represents the impactor. For reference, the value for the lowest velocities in our simulations is r grav, BH ≈ 10 −3 pc and r grav, NS ≈ 2 × 10 −4 pc. For b BH, NS ≲ r grav , we found alterations in both the inclination and eccentricity of virtually all retained planets regardless of the incidence angle and encounter velocity (see Appendices B and C ). The most efficient captures for both scenarios occurred in the range 60° ≲ θ BH, NS ≲ 120°, with up to eight planets for b BH, NS ≲ 10 −5 pc and the lowest impact velocity (see f cap in Figs. 2 and 3 ). Eccentricities of the captured planetary orbits, e cap , were very high – around 0.8 for the BH and 0.65 for the NS. Fig. 3. Outcomes of the NS scenario for v NS = 10 km s −1 . For each impact parameter (coded by colours and symbols) and incidence angle (columns), we plotted the final distribution of semi-major axes, eccentricities, and inclinations of planets that were either retained by the Sun or captured by the NS ( top six rows , for each planet from Mercury ☿ to Neptune ♆). Bottom two rows : fraction of retained and captured planets. Each variation in the initial parameters was averaged over nine pre-evolved states of the SS, see Sect. 2 . Detailed plots for all v NS and θ NS are in Figs. C.1 – C.8 . There was a difference between the prograde ( θ BH, NS = 0°) and retrograde ( θ BH, NS = 180°) encounters with the SS in both scenarios. A low-velocity NS showed no directional preferences for planetary capture, but it was most destructive for a retrograde fly-by with b NS ≲ 10 −5 pc; in 60% of the cases, the SS did not survive and in the remaining 40% the Sun was left with only one terrestrial planet (see f ret and f cap in Fig. 3 ). In contrast, a low-velocity BH was generally more disruptive and also more efficient in capturing planets for a prograde encounter (see f ret and f cap in Fig. 2 ). The greatest difference between the NS and BH case was found in the change of the orbital inclinations of the retained planets, i ret , as the BH is more likely to completely flip the orbits of even massive planets. Based on the positions of the planets during the encounter, both intruders were able to capture planets on prograde, retrograde, or 90° orbits while approaching from θ BH, NS = 0° (see i cap ). It was, however, impossible for the NS to capture a retrograde planet
Paper: 2019ApJ...872...53M (acknowledgments)
T.M. and E.R. are supported by DOE grant DE-SC0015975. E.R. acknowledges additional support by the Sloan Foundation, grant FG-2016-6443. J.L.T. and R.H.W. acknowledge support of NSF grant AST-1211889. J.D., R.H.W., S.M., and M.R.B. received support from the U.S. Department of Energy under contract number DE-AC02-76SF00515. Y.Y.M. is supported by the Samuel P. Langley PITT PACC Postdoctoral Fellowship. This research used resources of the National Energy Research Scientific Computing Center, a DOE Office of Science User Facility supported by the Office of Science of the U.S. Department of Energy under Contract No. DE-AC02-05CH11231. Software: Python, Matplotlib (Hunter 2007 ), NumPy (van der Walt et al. 2011 ), SciPy (Jones et al. 2001 ), emcee (Foreman-Mackey et al. 2013 ), george (Ambikasaran et al. 2015 ), CAMB (Lewis Bridle 2002 ), CLASS (Lesgourgues 2011 ), Rockstar (Behroozi et al. 2013 ), GADGET (Springel 2005 ), 2LPT (Crocce et al. 2006 ).
Paper: 2019A&A...628A..16R (fulltext)
in gas phase is still debated. In dark cores, only about 0.05% is locked in the S-bearing molecules detected at millimeter wavelengths (see the fractional abundances derived for S-bearing species towards the dark cloud TMC1-CP in Table 5 ). Unfortunately, most of the plausible major sulphur reservoirs such as S, HS, and S 2 are either not observable or very difficult to observe at millimeter wavelengths and one needs to trust to chemical models to estimate the exact amount of sulphur in the gas phase. It is widely accepted that most of the sulphur is locked in icy grain mantles. This is corroborated by the observations of S-bearing species in comets. In a recent compilation, Bockelée-Morvan Biver (2017) estimated that the S/O elemental ratio in the comet ice is about 1.5%, which is consistent within a factor of approximately four with the solar value, H 2 S being the main sulphur reservoir. Hot cores are compact warm regions where the icy mantles evaporate releasing the sulphur budget to the gas. Orion KL is the best-studied hot core, with full spectral surveys at millimeter and submillimeter wavelengths. In Table 5 , we show the abundances of sulphur species that have been derived in this prototypical hot core. Gas phase H 2 S is the main sulphur reservoir in this hot core and the elemental gas-phase sulphur abundance is ~ 1.6 × 10 −6 ( Crockett et al. 2014 ). Fuente et al. (2019) obtained a similar value for the sulphur gas phase abundance of ~10 −6 in the edges of the TMC1 molecular cloud. Table 5 Comparison of fractional abundances with observational templates. References. TMC1: Agúndez Wakelam (2013) ; Cernicharo et al. (2018) and references therein. Orion-KL: Tercero et al. (2010) ; Esplugues et al. (2014) ; Crockett et al. (2014) . Comets: Calmonte et al. (2016) ; Bockelée-Morvan Biver (2017) . Fig. 5 Abundances of NS + as a functionof N (H 2 ). Black pointscorresponds to the detections reported by Cernicharo et al. (2018) . Red points corresponds to the core and PDR positions in the Horsehead. The black solid line depicts a linear fit to the data. Based on the WHISPER survey, we have determined a significant inventory of sulphur species and their abundances in the Horsehead PDR. Only one H 2 S line has been observed and the estimated value of the H 2 S abundance (~0.3−5 × 10 −9 ) is hence uncertain. It seems clear, however, that the abundance of gaseous H 2 S is far from that observed in Orion-KL. One possibility is that a large fraction of H 2 S remains in the solid state. This is expected since the dust temperature in this PDR is around ~20–30 K in the PDR ( Goicoechea et al. 2009 ), which is below the sublimation temperature of H 2 S, ~80 K ( Jiménez-Escobar Muñoz Caro 2011 ). The detections of SO + and NS + suggest that a significant fraction of sulphur is in the form of S + in the outer layers of this cloud. Sulphur chemistry is still poorly known, and especially surface chemistry reactions need further understanding. Oba et al.
Paper: 2017A&A...600A..68G (acknowledgments)
We are grateful to our colleague P. Briand from the laboratory of mathematics LAMA of Université de Savoie Mont-Blanc for his kind help and for having provided us with a proof of the generalized central limit theorem. We also thank P. Mertsch who contributed by drawing our attention to the role of stable laws in cosmic ray physics.
Paper: 2017ApJ...834...15H (acknowledgments)
We thank Jo Bovy and Timothy D. Brandt for insightful discussions on statistical modeling of our results. C.P. acknowledges financial support from the Natural Sciences and Engineering Research Council of Canada. D.J.S. is supported by NSF grants AST-1412504 and AST-1517649. 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 Web Site is http://www.sdss.org/ . The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington. Based in part on observations obtained with MegaPrime/MegaCam, a joint project of CFHT and CEA/IRFU, at the Canada–France–Hawaii Telescope (CFHT) which is operated by the National Research Council (NRC) of Canada, the Institut National des Science de l’Univers of the Centre National de la Recherche Scientifique (CNRS) of France, and the University of Hawaii. This work is based in part on data products produced at Terapix available at the Canadian Astronomy Data Centre as part of the Canada–France–Hawaii Telescope Legacy Survey, a collaborative project of NRC and CNRS.
Paper: 2017ApJ...835...77M (acknowledgments)
We acknowledge the support from the ERC Consolidator Grant funding scheme ( project STARKEY , G.A. n. 615604). L.G. and T.S.R. acknowledge partial support from PRIN INAF 2014—CRA 1.05.01.94.05.
Paper: 2019A&A...625A.149J (fulltext)
by each successive refinement splitting a cell into two, and a refinement level of RL leads to cell sizes decreased by 1/2 RL . The maximum value of RL is taken as 13. This can lead to very small cells, with a commensurate decrease in the time step required for numerical stability. However individual loops can be simulated for reasonable real times (a few hours). Further details of the HYDRAD numerical method, including the finite difference schemes used, can be found in Appendix A.2 of BC13 and references therein. Such a “brute force” method is unlikely to be a viable way of running multi-dimensional codes. To this end we have developed an alternative approach, tested on one-dimensional problems ( Johnston et al. 2017a , b ), that treats the lower TR as an unresolved layer. By integrating the energy equation across this layer, and imposing a closure condition, we are able to provide rapid solutions to 1D problems with an accuracy that compares well with HYDRAD results. The details are found in Johnston et al. (2017a) , and we refer to this as the UTR method. It is incorporated into a one-dimensional version of the Lagrangian remap (Lare) code whose computational details are discussed in Arber et al. (2001) , referred to as “LareJ”. 2.2. Influence of numerical resolution We start by exploring the effect of numerical resolution on TNE cycles in coronal loop models. We model a coronal loop of total length 180 Mm with a small chromosphere attached to each end and select the largest grid cells in our calculations to have a width of 1 Mm. Thus, at the highest refinement level, the minimum cell width is 122 m. 2.2.1. Steady footpoint heating: HYDRAD simulations Representative of the conditions necessary to induce TNE in coronal loop models (e.g. Müller et al. 2003 ; Antolin et al. 2010 ; Peter et al. 2012 ; Mikić et al. 2013 ; Froment et al. 2018 ), we first consider the case of steady footpoint heating where the spatial profile is given by the sum of two Gaussian peaks (one at each loop leg), each defined as, These peaks are localised between the base of the corona and base of the TR with a maximal value at z = 12.5 Mm and we take z H = 1.5 Mm as the length scale of heat deposition. This is shown in Fig. 1 as the blue curve. We note that in this part of the paper a small spatially uniform background heating term is always present so that Q ( z ) = Q bg + Q H ( z ). This is commonly done in 1D loop models to ensure that T and n remain positive: here Q bg = 6.8682 × 10 −6 J m −3 s −1 and the effect of including a range of values of Q bg will be examined in Sect. 2.4 . Moreover, in order to avoid unrealistic pile up of condensations at the loop apex due to the symmetry of the model, the spatial symmetry of the heating profile is perturbed by adding a small enhancement of 0.4% to the Gaussian peak at the left-hand leg of the loop. The initial state of the loop is determined using just Q bg , leading to a temperature of order 1 MK. The footpoint
Paper: 2016MNRAS.462.3937L (acknowledgments)
Based on observations carried out as part of the Large Program ASAI (project number 012-12) with the IRAM 30 m telescope. IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain). This work was supported by the CNRS program ‘Physique et Chimie du Milieu Interstellaire’ (PCMI) and by a grant from LabeX Osug@2020 (Investissements d'avenir – ANR10LABX56). EM acknowledges support from the Brazilian agency FAPESP (grant 2014/22095-6 and 2015/22254-0 ). IJ-S acknowledges the financial support received from the STFC through an Ernest Rutherford Fellowship (proposal number ST/L004801/1 ).