adsabs/WIESP2022-NER huggingface dataset:
Paper: 2019ApJ...886..114H (fulltext)
of both the non-MAD and MAD states. Our main conclusions are summarized as follows: 1. High-energy particles at the first (strongest) shock during the debris circularization phase are unlikely to be produced, because the mass fallback rate exceeding the Eddington rate makes both the Compton drag and the proton–proton collision cooling very efficient. Since the shocked region is far from the black hole, the magnetic field is relatively weak. The weak magnetic field is another reason whythe stochastic acceleration of the protons becomes inefficient. 2. High-energy particles during the super-Eddington accretion phase of the non-MAD state are unlikely to be produced because of very efficient proton–proton collision cooling and Compton drag, even though the stochastic acceleration of the protons is more efficient because of the much stronger magnetic field than that of the first shock case. 3. Neutrinos and gamma-ray photons are produced by pions not only in the super-Eddington MAD scenario but also in the RIAF scenarios of both the non-MAD and MAD cases. However, the gamma-ray photons cannot escape from the source because of the large optical depth in the super-Eddington MAD state and because of the very efficient pair production in the RIAF cases. 4. In the RIAF phase, the protons can be accelerated up to with the other given appropriate parameters. The neutrino energy emitted at the peak of the spectrum is then estimated as . The neutrino emission is proportional to t −10/3 , which is steeper than the standard TDE decay rate. This indicates that one can potentially identify whether an RIAF that existed around an SMBH has a TDE origin or not. 5. In the super-Eddington MAD state, the stronger magnetic field than in the non-MAD case makes it possible to accelerate the protons up to ultrarelativistic energies. The dominant process to prevent the protons from accelerating strongly depends on the black hole mass. If the black hole mass is larger than , synchrotron cooling is the most dominant process. In this case, the protons are accelerated up to an energy of with the other given appropriate parameters. The neutrino energy at the peak of the spectrum is then estimated as . Otherwise, the Compton drag is more efficient process. In this case, the resultant energy of the protons increases at most up to . 6. In the super-Eddington MAD state, if M bh ≳ 10 7.7 M ⊙ , the neutrino light curve is proportional to , which is steeper than the standard TDE decay rate. On the other hand, it follows the decay rate with energy more than if . In both cases, the neutrino luminosity is of the order of the Eddington limit or the super-Eddington limit. Such a high neutrino luminosity and characteristic light curve give us the ability to judge whether the TDE disk is in the MAD state or not. 7. In the radiatively inefficient MAD state, the stronger magnetic field than in the non-MAD case makes it possible to accelerate the protons up to . The neutrino energy estimated
Paper: 2016A&A...585A.140L (acknowledgments)
The Swedish 1-m Solar Telescope is operated on the island of La Palma by the Institute for Solar Physics of Stockholm University in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofísica de Canarias. The author is grateful to Mikael Ingemyr and Pit Sütterlin for help with the observations, to Guus Sliepen for making it possible to do the DA scans, and to Göran Scharmer for discussions and helpful suggestions.
Paper: 2018A&A...617L...6R (fulltext)
He atom, which includes 4 levels for He II and 1 level for He III, as well as a 5-level-plus continuum H atom. By calculating the full non-LTE radiative transfer for each model we find the energy level populations and thus ionization densities for both atomic species. The cylinder is orientated such that its axis is parallel to the solar surface. We then calculate the brightness temperature for a set of horizontal lines-of-sight through the prominence cylinder. We follow the same methods for brightness temperature calculation as outlined in Rodger Labrosse (2017) with one notable change in the calculation of the thermal gaunt factor. In this study we no longer estimate the thermal gaunt factor through the assumption given in Dulk (1985) for plasma of temperature 2 × 10 5 K, but instead take our values by interpolating the table of calculated thermal gaunt factors of van Hoof et al. (2014) , as described in Gayet (1970) and Simões et al. (2017) . Whilst the optical thickness diagnostic discussed in this study assumes purely thermal bremsstrahlung absorption, the results from these numerical models are calculated using both thermal bremsstrahlung and neutral hydrogen absorption. For this analysis we first consider a set of isothermal prominence models to see how well they agree with Eq. ( 4 ) and see how this changes when considering a multi-thermal plasma instead. Whilst nominally the models used here describe solar prominences, the results are applicable to any off-limb solar atmospheric structure. For on-disc structures, the spectral gradient will follow a similar relation to Eq. ( 4 ) with an additional term corresponding to the contribution of the background solar continuum spectrum. On-disc structures would therefore require knowledge of the brightness temperature of the structure and the background emission illuminating it from the solar disc. This will likely be problematic unless the structure in observation is formed above the chromosphere, the forming region for most millimetre radiation, or where the structure is transient in nature such that measurements of both the background and enhanced brightness temperature phases are obtainable. 3.1. Isothermal models The set of isothermal models we use here are the same as the “ t ” models as described in Gouttebroze Labrosse (2009) . They each consider a prominence cylindrical cross-section of radius 0.5 Mm at an altitude of 10 Mm. The isobaric pressure is 0.1 dyn cm −2 , the helium abundance is 0.1 and the micro-turbulent velocity is 5 km s −1 in all models. The set of temperatures considered is {6000, 8000, 10 000, 15 000, 20 000, 30 000, 40 000, 50 000, 65 000, 80 000, 100 000} K. From Fig. 1 it can be seen that whilst the relationship between the modelled spectral gradient and the optical thickness follows a similar trend to that expected by Eq. ( 4 ) the values are slightly lower. This discrepancy is caused by the assumption that the gaunt factor is approximately constant over the
Paper: 2019MNRAS.482.2039C (fulltext)
that R 90 is highly dependent on the finding algorithm. Interacting galaxies show a very wide f Y distribution for all quantities studied. There is an evident difference between host and satellite galaxies, which peak at f Y > 1 and f Y 1, respectively. These differences are mainly caused by inadequate particle assignment in halomaker , which we show our improved version of velociraptor handles better. halomaker artificially increases (decreases) the estimated values of M * , SFR, and R 50,90 for host (satellite) galaxies in interacting systems. Differences between the catalogues are amplified at higher redshifts, where the f Y distributions of interacting and isolated galaxies widen. Galaxy population statistics . We investigate how the choice of finder affects the overall galaxy population statistics. We explore the GSMF, SFRF, and mass–size relation for R 50 and R 90 . At M * 10 11 M ⊙ the z = 0 GSMFs of halomaker and velociraptor agree well, while at higher stellar masses the former predicts from 20 per cent to 100 per cent more galaxies than the latter. At higher redshifts differences are amplified. At z = 4 halomaker ’s GSMF predicts a number density of galaxies at least 30 per cent higher than that of velociraptor over the whole mass range, increasing to 100 per cent at M * ≳ 10 10 M ⊙ . The SFRF at z = 0 is also fairly similar between the finders, with differences increasing with redshift. At z = 2, the peak of the cosmic star formation history, halomaker predicts up to 50 per cent more galaxies at than velociraptor . This is important as these galaxies are expected to dominate the cosmic SFR. We compare the R 50 and R 90 size–mass relation predicted by both finders. We find that the R 50 mass–size relations resulting from the two finders are similar, except at the high-mass end, M * ≃ 10 12 M ⊙ , where halomaker’s galaxies are 20 per cent larger than velociraptor’s . These differences increase by 30 per cent when we study R 90 . This results from the fact that the stellar content and structure in the outskirts of galaxies are very sensitive to the choice of finder. Although we see that the overall z = 0 galaxy statistics are not greatly impacted by the choice of finder, individual galaxies can display differences in mass and size of more than a factor of 3 between the two finders studied here. We suggest that the tuning of simulations of galaxy formation is relatively robust as it has consistently focused on population statistics. However, comparisons of galaxy subpopulations with observations, specifically in the context of pairs, groups, and clusters, can be greatly affected by the choice of finder. We showed that our new algorithm outperforms 3D finders and provides extensive evidence of this. One of our key findings is that the stellar outskirts of galaxies is greatly affected by the choice of finder. In upcoming studies we will explore in detail the diffuse IHSC, the stellar streams, and the outer stellar profiles of
Paper: 2016MNRAS.459.4029L (acknowledgments)
This research was supported in part by the Natural Science Foundation of China under Grants Nos 11175270 and 11475065, and the Program for New Century Excellent Talents in University under Grant No. NCET-12-0205.
Paper: 2016ApJ...832...39L (acknowledgments)
We thank the anonymous referee for helpful comments that greatly improved the presentation of this paper. This work is supported by 973 Program (No. 2015CB857002), National Science Foundation of China (grant Nos. 11203054, 11128306, 11121062, 11233005, and 11073017), NCET-11-0879, the Strategic Priority Research Program “The Emergence of Cosmological Structures” of the Chinese Academy of Sciences, grant No. XDB09000000 and the Shanghai Committee of Science and Technology, China (grant No. 12ZR1452800). We also thank the support of a key laboratory grant from the Office of Science and Technology, Shanghai Municipal Government (No. 11DZ2260700). H.J.M. would like to acknowledge the support of NSF AST-1517528. A computing facility award on the PI cluster at Shanghai Jiao Tong University is acknowledged. This work is also supported by the High Performance Computing Resource in the Core Facility for Advanced Research Computing at the Shanghai Astronomical Observatory.
Paper: 2016MNRAS.463.4317P (acknowledgments)
We thank the anonymous reviewer for their helpful comments, which improved the content of this paper. AHP and SZ thank the Lady Davis Foundation and The Netherlands Organization for Scientific Research (NWO) VICI grant for the financial support. LVEK and BKG acknowledge the financial support from the European Research Council under ERC-Starting Grant FIRSTLIGHT – 258942. AGdB, SY, MM and VNP acknowledge support by the ERC for project 339743 (LOFARCORE). VJ acknowledges the NWO for the financial support under VENI grant – 639.041.336 . ITI was supported by the Science and Technology Facilities Council [grant number ST/L000652/1 ]. The LOFAR was designed and constructed by ASTRON, the Netherlands Institute for Radio Astronomy, and has facilities in several countries, which are owned by various parties (each with their own funding sources) and are collectively operated by the International LOFAR Telescope (ILT) foundation under a joint scientific policy.
Paper: 2018ApJS..237...33X (fulltext)
when the Chabrier IMF is adopted. However, we argue that a 20% contribution from the thick disk at the local disk is seriously overestimated. Our results suggest that the thick disk contributes only a few percent mass density at the disk midplane, which is consistent with many literature results (e.g., Jurić et al. 2008 ; Chen et al. 2017 ). We emphasize that our results are obtained at solar radius but not the “solar neighborhood.” Our sample stars have a good coverage at 400–800 pc away from the Sun in the azimuthal direction but have poor coverage within 400 pc. A likely explanation of the higher density found by this work than the solar-neighborhood values in the literature is that the Sun is located in a local low stellar density region, which has a density of 0.0164 M ⊙ pc −3 (or 0.0106 M ⊙ pc −3 if the Chabrier IMF adopted) lower than the nearby disk. Such a difference must be contributed by the Local stellar arm. Our Sun is either located at the inner boundary of the Local arm or embedded in a cavity of stars in the arm, and it needs to be further studied using data with improved spatial coverage to clarify which is the real case. Note that the literature results for the solar-neighborhood density are usually determined within a complex volume, which varies with different types of stars. It is thus difficult to make a direct comparison of our relatively well defined volume density with the literature results. To test whether the difference is caused by the possibility that the literature results are actually averaged values in a larger volume, we have also examined the mean stellar mass density within 7.8 kpc R 8.2 kpc and pc, and we find a density of 0.0549 M ⊙ pc −3 (0.0496 M ⊙ pc −3 from the Chabrier IMF), which is still significantly higher than the “solar-neighborhood” values in the literature. We also emphasize that since the stellar mass density decreases quickly with increasing height above the disk plane, the “underlying” midplane density should be higher than the current estimates of average values within pc. The midplane density is expected to be comparable to the measured values without correction for contaminations of main-sequence stars (Section 4 ). Assuming a gas density of 0.05 M ⊙ pc −3 as widely adopted (Holmberg Flynn 2000 ; Flynn et al. 2006 ), the expected mass density of baryon matter (star and gas) in the nearby disk plane within a few hundred parsecs is thus 0.109 M ⊙ pc −3 (0.104 M ⊙ pc −3 for the Chabrier IMF). Such a value is highly consistent with the local total mass density yielded by stellar dynamics, which suggest a typical value of 0.1 M ⊙ pc −3 , with an uncertainty of ∼0.01 M ⊙ pc −3 (Bienayme et al. 1987 ; Kuijken Gilmore 1989c ; Pham 1997 ; Holmberg Flynn 2000 ; Read 2014 ; McKee et al. 2015 ; Widmark Monari 2017 ; Kipper et al. 2018 ). Our results thus leave little room for the existence of a meaningful amount of dark matter in the nearby disk midplane. However, since our results show that
Paper: 2019MNRAS.488..525Z (acknowledgments)
ACKNOWLEDGEMENTS This work received generous support from the National Natural Science Foundation of China, project nos 11763007, 11863005, 11803026 and 11503008. We would also like to express our gratitude to the Tianshan Youth Project of Xinjiang, no. 2017Q014.
Paper: 2018ApJ...853...50F (acknowledgments)
We are grateful to the anonymous referee whose excellent comments and suggestions greatly improved this article. This work is supported by grants from the HKRGC (HKU7062/13P) and the Seed Fund for Basic Research at the HKU (201611159037). This work is partly supported by grants DGAPA/PAPIIT-107215 and CONACyT-CB2015-254132. R.G.B. acknowledges support from the Spanish Ministerio de Economía y Competitividad, through projects AYA2016-77846-P and AYA2014-57490-P. X.W.L. acknowledges financial support from the National Key Basic Research Program of China 2014CB845700. M3.B. is partially supported by ANPCyT through grant PICT-2014-2708 and by MinCyT-DAAD bilateral cooperation through grant DA/16/07. X.F. is grateful to Albert A. Zijlstra, Bruce Balick, and Romano L.M. Corradi for in-depth discussion. We thank Paolo Ventura for providing us the AGB model predictions. We also thank Michael R. Merrifield for giving us the permission to reproduce the stellar orbit (in Figure 2 of this paper) based on Figure 2 in Merrett et al. ( 2003 ). This research also made use of NASA’s Astrophysics Data System ( http://adsabs.harvard.edu ), the SDSS DR10 Science Archive Server (SAS; http://data.sdss3.org/ ), the SDSS DR12 Finding Chart Tool ( https://skyserver.sdss.org/dr12/en/tools/chart/ ), and the SIMBAD Astronomical Database ( http://simbad.u-strasbg.fr/simbad/ ). Software: iraf , midas , SAOImage DS9. Facility: GTC (OSIRIS).
Paper: 2017A&A...605A..87F (acknowledgments)
We are grateful to the University of Catania for providing the funds necessary to carry out the Observational Campaign in La Palma. The Swedish 1-m Solar Telescope is operated on the island of La Palma by the Institute for Solar Physics (ISP) of Stockholm University in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofísica de Canarias. We are also grateful to Professor Giovanni Gallo of the University of Catania for his contribution in the implementation of the new segmentation algorithm. This research has received funding from the EC 7th Framework Programme FP7/2007–2013 under the grant agreement SOLARNET (project No. 312495). This work was also supported by the Istituto Nazionale di Astrofisica (PRIN INAF 2014), and by the Italian Ministry of University and Research (PRIN MIUR 2012).
Paper: 2018ApJ...852...70B (acknowledgments)
We thank the US National Science Foundation (AST-1505502) for support in conducting the present experiments. The authors would like to acknowledge the W. M. Keck Foundation for financing the experimental setup.
Paper: 2019ApJ...871..172R (fulltext)
( 2015 ), Lestrade Thilliez ( 2015 ), Booth et al. ( 2017 ), Chavez-Dagostino et al. ( 2016 ), and this paper (see Table 2 ). In the case of the observed millimeter emission, the rapid rise of the flux density with frequency (a spectral index of 4.9 ± 1.8) favors emission from warm dust (e.g., Ricci et al. 2010 ). In contrast, the centimeter emission is remarkably flat (a spectral index of −0.07 ± 0.25), suggesting an optically thin free–free nature. Below we examine the expected properties of the free–free emission. Figure 2. Spectrum of the emission associated with Eridani at centimeter and millimeter wavelengths from the references given in the text. The long-dashed lines are separate power-law fits to the centimeter and to the millimeter emission. The data point at 3 GHz (0.48 in the logarithm) is a 3 σ upper limit and was not taken into account in the fit. 4.1. Optically Thin Free–Free Emission The free–free power per unit volume per Hz is given bywhere Z is the ion charge, n e and n i are the electron and ion number density, respectively, T is the plasma temperature, and is the velocity-averaged Gaunt factor. The total power per unit volume is obtained integrating over the frequency where is the frequency average of the velocity-average Gaunt factor, which is in the range 1.1–1.5. Choosing gives a 20% accuracy (Rybicki Lightman 1986 ). The power per unit volume over a frequency range is given bywhere the normalized energy is u = hν / kT . For an optically thin plasma at a constant temperature, the luminosity in a normalized energy range ( u 1 , u 2 ) is obtained by integrating the emissivity over the volume V : The X-ray luminosity of Eridani has been measured with XMM-Newton in the energy range [ E 1 , E 2 ] = [0.2, 2] kev, with values in the range (1.5–1.68) × 10 28 erg s −1 (Poppenhaeger et al. 2010 ; Loyd et al. 2016 ). Let L R be the luminosity in the frequency range [ ν 1 , ν 2 ] = [6, 43] GHz where the spectrum is flat. Then, from Equation ( 5 ), the ratio of the free–free luminosity in the X-ray range L X and in the radio range L R is given bywhere the lower value corresponds to T = 2 × 10 6 K and the upper value corresponds to T = 3 × 10 6 K. These temperatures are similar to those measured by Schmitt et al. ( 1996 ) for the peak of the temperature distribution of the coronal material of Eridani. The observed radio luminosity in the range 6–43 GHz is given bywhere D = 3.2 pc is the distance to Eridani. Then, this radio luminosity implies a free–free emission in X-rays offor a temperature in the range (2–3) × 10 6 K, in agreement with the observed X-ray luminosity observed by XMM-Newton . Therefore, the radio emission in the flat-spectrum region is consistent with free–free emission for this range of plasma temperature. Note that line X-ray emission is also important in hot plasmas (e.g., Mewe et al. 1985 ; see their Figure 1). Thus, the total X-ray emissivity should be larger than the continuum free–free emission calculated in
Paper: 2018ApJ...862...33A (fulltext)
lower-mass stars with which to anchor the comparison to the models. Still, it does appear that the model is significantly overpredicting the spin-down for the 0.6–0.8 M ⊙ stars, with the predicted median star in that mass range having a P rot = 21.0 days at 1.344 Gyr, ≈4.5 days more than what is observed. The four lower-mass NGC 752 members have rotation periods that are more consistent with what is predicted by the Matt et al. ( 2015 ) model, with the median 0.4–0.6 M ⊙ star predicted to have a P rot = 17.2 days. One possible interpretation is that we are seeing the evolutionary stalling observed in the comparison of Praesepe and NGC 6811 for 0.8–0.9 M ⊙ stars shifted to lower masses, with the 0.6–0.8 M ⊙ stars being the ones now rotating significantly faster than expected at this age. 6.2. Comparing Magnetic Activity in NGC 752 and in the Hyades and Praesepe Studies of observational tracers of coronal or chromospheric activity have uncovered a mass-dependent transition between active and inactive stars in open clusters (e.g., Kafka Honeycutt 2006 ; Douglas et al. 2014 ; Núñez Agüeros 2016 ; Núñez et al. 2017 ). The dividing line between these two populations shifts to lower masses in older clusters, indicating that lower-mass stars possess longer activity lifetimes. For FGK stars, these lifetimes are ≤650 Myr, as calibrated by observations of open clusters younger than NGC 752 (Hawley et al. 1999 ; Douglas et al. 2014 ; Núñez et al. 2017 ). Extending such measurements to older open clusters is a primary motivation of this work. Our knowledge of the chromospheric activity lifetimes of lower-mass stars currently relies on indirect calibrations, such as modeling the vertical gradient in H α emission line strengths as a consequence of dynamical heating in the Galactic disk (e.g., West et al. 2008 ). Our spectroscopic campaign confirms that the boundary between active and inactive stars has shifted well into the M dwarf regime in NGC 752. As noted above, there are three stars in our spectroscopic sample with formal detections of H α emission, but we do not consider these stars indicative of the location of the active/inactive boundary in this cluster. As Figure 11 demonstrates, even the modest activity signatures measured from these stars (EW >2 Å) make them clear outliers from the dominant cluster locus, and their moderate membership probabilities (50 P mem 80%) indicate that these stars may not be bona fide cluster members. Calculating the mean H α EW for NGC 752 members in bins of ( r − K ), as shown in Figure 15 , indicates that there is no transition to activity in NGC 752, at least within the domain of our spectroscopic survey, which includes stars with SpTs as late as ≈M2. Figure 15. Average H α EW as a function of logarithmically binned ( r − K ) for stars in NGC 752 compared to the younger Praesepe and Hyades clusters. The number of stars in each bin is along the top; the vertical bars show the σ for the bin, and the horizontal bars show