adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016MNRAS.463.2600L (acknowledgments)
We acknowledge the dedicated team behind the Kepler and K2 missions, without whom this work would not have been possible. We thank Daniel Huber and Benoit Mosser for useful comments on an earlier version of the paper. MNL acknowledges the support of The Danish Council for Independent Research | Natural Science (Grant DFF-4181-00415). MNL was partly supported by the European Community's Seventh Framework Programme (FP7/2007-2013) under grant agreement no. 312844 (SPACEINN), which is gratefully acknowledged. Funding for the Stellar Astrophysics Centre (SAC) is provided by The Danish National Research Foundation (Grant DNRF106). The research was supported by the ASTERISK project (ASTERoseismic Investigations with SONG and Kepler ) funded by the European Research Council (Grant agreement no.: 267864). WJC, GRD, and AM acknowledge the support of the UK Science and Technology Facilities Council (STFC). SB acknowledges partial support from NASA grant NNX13AE70G and NSF grant AST-1514676. AMS is partially supported by grants ESP2014-56003-R and ESP2015-66134-R (MINECO). VSA acknowledges support from VILLUM FONDEN (research grant 10118). RAG acknowledges the support from the ANR program IDEE (n. ANR-12-BS05-0008) and the CNES. DWL acknowledges partial support from the Kepler mission under Cooperative Agreement NNX13AB58B with the Smithsonian Astrophysical Observatory. This research has made use of the Washington Double Star Catalog maintained at the U.S. Naval Observatory; the WEBDA data base, operated at the Department of Theoretical Physics and Astrophysics of the Masaryk University; and the SIMBAD data base, operated at CDS, Strasbourg, France.
Paper: 2016A&A...589A.116M (acknowledgments)
This work has been financially supported by the Programme National Hautes Énergies (PNHE). The research leading to these results has received funding from the European Union Seventh Framework Program (FP7/2007-2013) under grant agreement No. 312789. The XMM-Newton project is an ESA Science Mission with instruments and contributions directly funded by ESA Member States and the USA (NASA). This work is based on observations made with the NASA/ESA Hubble Space Telescope obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. These HST observations are associated with programs 13403 and 13316. This work is based on observations made with ESO Telescopes at the Paranal Observatory under programs 091.B-0183(H), 092.B-0920(A) and 093.B-0932(A). Karl G. Jansky Very Large Array (VLA) of the National Radio Astronomy Observatory is a facility of the National Science Foundation, operated under a cooperative agreement by Associated Universities, Inc. Support for CARMA construction was derived from the states of California, Illinois, and Maryland, the James S. McDonnell Foundation, the Gordon and Betty Moore Foundation, the Kenneth T. and Eileen L. Norris Foundation, the University of Chicago, the Associates of the California Institute of Technology, and the National Science Foundation. Ongoing CARMA development and operations are supported by the National Science Foundation under a cooperative agreement, and by the CARMA partner universities.
Paper: 2017ApJ...851...77R (fulltext)
given the age determined by Burgasser Mamajek ( 2017 ), the activity index described in Section 3.2 , and the XUV fluxes that have been measured for the star (Bolmont et al. 2017 ; Bourrier et al. 2017 ; Wheatley et al. 2017 ). As noted earlier, the combination of these stellar attributes ensures that the TRAPPIST-1 planets will have been subjected to extended periods of activity-related atmospheric erosion. As such, the possible atmospheric erosion effects will need to be accounted for in consideration of the habitability potential of the planets. We calculated the incident flux and escape velocities for the TRAPPIST-1 planets based on the planetary and stellar properties provided by Gillon et al. ( 2017 ), where the bolometric luminosity is the result of spectral energy distribution analysis by Filippazzo et al. ( 2015 ). These calculated values are shown in Table 1 , and are depicted in Figure 5 along with the incident flux and escape velocities for various solar system objects. The cosmic shoreline based upon the estimates by Zahnle Catling ( 2017 ) is shown as a solid line, where the solar system objects to the left of the line tend to not have atmospheres and objects to the right tend to have (albeit sometimes tenuous) atmospheres. Note that the incident flux shown in the plot represents the total (bolometric) flux received at the distance of the planets from the host star. The derivation of the cosmic shoreline power law by Zahnle Catling ( 2017 ) utilized the scaling relationships of Lammer et al. ( 2009 ) that related extreme ultraviolet (EUV) and XUV excess to stellar age and spectral type. Thus, our adoption of the bolometric flux includes the atmospheric erosion effects of the intrinsic EUV and XUV radiation as a subset of the total flux. The XUV flux related to magnetic activity, as described in this paper, will add an additional component to the atmospheric eroding flux effects. Increasing the XUV component has the effect of down-weighting the influence of the escape velocity for retaining the atmosphere. This means that low-mass stars with a higher relative rate of XUV emission will have a substantially greater atmospheric erosion effect on their hosted planets. In Figure 5 , we include two additional lines for the cosmic shoreline that reduce the escape velocity effect by factors of 2 (dashed) and 4 (dotted). These lines thus represent stars for which the XUV flux is 16 and 256 times higher than the solar flux. In the case of the former, this indicates that only the g planet has a reasonable chance of retaining a planetary atmosphere, whereas in the case of the latter, all of the planets are likely to be barren. These shifts of the cosmic shoreline are consistent with the XUV flux measurements conducted by Wheatley et al. ( 2017 ) of TRAPPIST-1, which indicated that the incident flux received by the planets may be several orders of magnitude larger than that received by the Earth. Such incident flux would result in atmosphere
Paper: 2015ApJ...808...28C (acknowledgments)
We thank Anvar Shukurov and Andrew Fletcher for valuable discussions and for helpful comments on an early draft of the manuscript, and David Moss for useful feedback on a later version. We are also grateful to Romeel Davé Oliver Gressel, Marita Krause, Jonathan Zwart, Nathan Deg, Prasun Dutta, Pallavi Bhat, and especially Kandaswamy Subramanian for useful discussions. Finally, we thank the anonymous referee for insightful suggestions that helped to improve the paper.
Paper: 2016A&A...585A.106B (fulltext)
flow direction forms with the grain’s velocity. The first part of Eq. ( 9 ) is the sum of radiation pressure force with the corpuscular pressure force split into its radial and transverse component. The latter part of Eq. ( 9 ) is the sum of the classical PR drag with the solar wind drag ( Marzari Vanzani 1994 ). The adopted numerical algorithm for solving the equations of motion is the RA15 version of the RADAU integrator ( Everhart 1985 ). 3. Initial conditions for asteroidal dust and cometary dust particles To compute the flux of dust particles released in the main belt through asteroidal collisions and reaching the Earth and Mercury, we adopted an initial distribution of orbital elements with the following characteristics ( Borin et al. 2009 ; Marzari et al. 1996 ; Dohnanyi 1969 ). The initial semimajor axis is randomly selected in between 2.1 and 3.3 AU, the initial eccentricity varies in the range 0.00.4 and the inclination is in the range . This choice reflects the typical orbital elements of the asteroid belt ( Gradie et al. 1989 ; Zappalá Cellino 1993 ; Zappalá et al. 1994 ; Milani Knezevic 1994 ). We assume a grain density of , which is a reasonable value for dust particles coming from the main belt asteroids ( Grün et al. 1985 ). Spherical particles are considered in the approximation of Mie’s theory for which , the dimensionless, radiation-pressure coefficient averaged over the solar spectrum, is 0.53 ( Marzari Vanzani 1994 ; Mukai et al. 1982 ). For the JFC dust particles we considered orbits with an initial semimajor axis in a range between 2.5 and 26.7 AU with the majority of dust grains in the range between 2.5 and 8 AU, eccentricity in a range interval between 0.3 and 0.9, and orbit inclination in the range between with a large concentration of particles in the range between . In this case, we assume a density of , the radiation pressure coefficient following the asteroids grain simulations ( Marzari Vanzani 1994 ; Mukai et al. 1982 ). The initial orbital elements of comets are taken from Nesvorny et al. (2010) and Levison Duncan (1997) , who followed the evolution of bodies originating in the Kuiper Belt as they are scattered by planets and evolve into the inner solar system. In their input list of source objects, Nesvorny et al. (2010) consider a JFC to become active at the time when the perihelion distance, q first drops below 2.5 AU in the Levison Duncan (1997) simulations. The orbital distribution of visible JFCs obtained with this model by Levison Duncan (1997) well approximates the observed distribution. Defining as the fading lifetime of JFCs, which is the time range between their first and last appearance ( Levison Duncan 1997 ; Nesvorny et al. 2010 ), we considered both the initial time and the final time () for each comet with perihelion distance . However, the outcome in terms of close encounters and impact speed was not significantly different, so we used the set with the final time. In Fig. 1 we illustrate the
Paper: 2016A&A...585A.148B (acknowledgments)
We thank the referee for his/her accurate revision of the paper. B.B. acknowledges support from grant PRIN-INAF 2011 “Black hole growth and AGN feedback through the cosmic time”. M.B. acknowledges support from the FP7 Career Integration Grant “eEASy”: “supermasssive black holes through cosmic time: from current surveys to eROSITA-Euclid Synergies” (CIG 321913). Funding for SDSS-III has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, and the US Department of Energy Office of Science. The SDSS-III web site is http://www.sdss3.org/ . SDSS-III is managed by the Astrophysical Research Consortium for the Participating Institutions of the SDSS-III Collaboration including the University of Arizona, the Brazilian Participation Group, Brookhaven National Laboratory, Carnegie Mellon University, University of Florida, the French Participation Group, the German Participation Group, Harvard University, the Instituto de Astrofisica de Canarias, the Michigan State/Notre Dame/JINA Participation Group, Johns Hopkins University, Lawrence Berkeley National Laboratory, Max Planck Institute for Astrophysics, Max Planck Institute for Extraterrestrial Physics, New Mexico State University, New York University, Ohio State University, Pennsylvania State University, University of Portsmouth, Princeton University, the Spanish Participation Group, University of Tokyo, University of Utah, Vanderbilt University, University of Virginia, University of Washington, and Yale University.
Paper: 2020A&A...635A..52M (fulltext)
with 10% and 30% of binaries. In the latter ones, the peak density is around the (0,0) point. The binaries contribute to producing an elongated structure with a lower density (which increases with the binary fraction). In NGC 5272 3 , the P1 sequence is relatively homogeneously populated with barely a very small overdensity near the origin. The same trend is observed for NGC 6254 (not shown), from which we conclude that either the binary fractions in GCs are underestimated, or binaries contribute only part of the extension of P1, and another process is at play. Fig. 7. Hess diagrams based on chromosome maps of the synthetic clusters with 10% ( top left panel ) and 30% ( top right panel ) of binaries, together with that of of NGC 5272 ( bottom panel ). A further test of the effects of binaries on the chromosome map is shown in Fig. 8 . We selected the cluster NGC 6254 since; (1) it shows an extended P1 ( Milone et al. 2017 ), (2) its metallicity corresponds to that of our models, and (3) the binary fraction decreases from the core to the outer parts of the cluster ( Dalessandro et al. 2011 ; Milone et al. 2012 ). To investigate whether or not this variation of the binary fraction is reflected in the shape of P1, we selected two regions of the cluster: the core and an annulus around it, as displayed in the left panel of Fig. 8 . We subsequently selected the stars brighter than the bottom of the RGB in each region and constructed their chromosome map. The results, shown in the right panel of Fig. 8 , do not show any significant variation of the extent of the chromosome map. This result does not depend on the choice of the core and annulus regions. If binaries governed the elongation of P1, one would expect a smaller extension in the outer regions of the cluster, which is not observed. Fig. 8. Left panel : position of stars brighter than m 814 W = 18.5 in the cluster NGC 6254. The red and blue circles indicate the two regions used to produce the chromosome maps of the right panel . Right panel : chromosome map of the stars inside the blue circle in the left panel (blue squares) and of the stars in between the red and blue circles (open red circles). Only stars brighter than m 814 W = 16.5 (i.e., above the bottom of the RGB) are considered in the chromosome maps. If binarity were the main reason for the extension of P1, one would also find a correlation between binary fraction and P1 extension. A simple comparison of two GCs with similar properties, NGC 5272 (M 3) and NGC 6205 (M 13), shows some qualitative trends. Both clusters are of about the same age, metallicity and mass ( Harris 1996 ; Marín-Franch et al. 2009 ), but M 3 shows an extended P1, while M 13 does not – see Fig. 5 of Milone et al. (2017) . According to Milone et al. (2012) , the total binary fraction in M 3 (M 13) is between 3 and 6% (2 and 10%). Thus, M 3 does not have a higher binary fraction that would explain its more extended P1. 3.3. Effect of the binary mass ratio
Paper: 2018ApJ...853..179T (acknowledgments)
We are grateful to the staff of the IRAM facilities for the continuing excellent support of the large CO survey programs we have been analyzing in this paper. We also thank Albrecht Poglitsch and Matt Griffin, the PACS and SPIRE teams, and ESA for their excellent work on the Herschel instruments and mission. Finally, we thank the referee, whose suggestions have led to a clearer manuscript overall. This paper makes use of the following ALMA data: ADS/JAO.ALMA#2013.1.00092.S. ALMA is a partnership of ESO (representing its member states), NSF (USA), and NINS (Japan), together with NRC (Canada), MOST and ASIAA (Taiwan), and KASI (Republic of Korea), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO, and NAOJ.
Paper: 2020ApJ...890...55L (fulltext)
unavoidable; that is, first-principles numerical simulations that address this plasma regime with adequate separation between the ion and the electron scales may inevitably satisfy the conditions for nonlinear reconnection in turbulent eddies at sub-ion scales. Indeed, the biggest such simulation that we are familiar with reports spectra steeper than in the sub-ion range (Told et al. 2015 ), which could be a manifestation of reconnection-driven energy depletion in the eddies. It is interesting to analyze this result in light of recent MMS observations of so-called electron-only reconnection in the Earth's magnetosheath (Phan et al. 2018 ; Stawarz et al. 2019 ). In Boldyrev Loureiro ( 2019 ) we have estimated that the decoupling of the ions in these events requires or depending on whether one assumes a or magnetic field profile for the reconnecting field in the eddy. These estimates range from ∼6 to ∼12; combined with Equation ( 12 ) or ( 13 ) they suggest a rather narrow range of scales where nonlinear “electron-only” reconnection in the eddies may be possible. Remarkably, Phan et al. ( 2018 ) report a current sheet thickness of , strikingly consistent with these numbers and with Equation ( 12 ). This is certainly very encouraging, but one must also bear in mind that all our analytical results are only order of magnitude estimates which ignore order unity numerical prefactors. Another observationally based result which we interpret to be consistent with our analysis is the recent claim by Chen et al. ( 2019 ) that energy dissipation at kinetic scales in the magnetosheath is dominated by linear electron Landau damping (the energy dissipation rate via that channel being comparable to the energy cascade rate). Indeed, Equation ( 12 ) demonstrates that full reconnection in sub-ion scale eddies is permitted for typical magnetosheath parameters; and previous investigations of heating in (strong guide-field) collisionless reconnection (Loureiro et al. 2013 ; Numata Loureiro 2015 ) show that when β e ≪ 1 linear electron Landau damping is by far the dominant energy dissipation channel. Furthermore, Chen et al. ( 2019 ) report a steepening of the spectrum at the scales where Landau damping is prominent, consistent with our expectations of the effects of active collisionless reconnection in the cascade. 4.3. Low-beta Pair Plasmas For completeness, we address here the case of low-beta pair plasmas, for which governing equations and turbulence scalings are given in Loureiro Boldyrev ( 2018 ). In the MHD regime the same equations as for an electron-ion plasma apply; the discussion is thus unchanged from that of Section 3 . In the case of a collisionless pair plasma, the criterion stated in Equation ( 3 ) yields (exactly as in Equation ( 12 ) since the scaling for γ t is the same). This needs to be reconciled with the requirements for the existence of a tearing range in the first place: for a Harris-type reconnecting profile, these are (Loureiro Boldyrev
Paper: 2018MNRAS.474.3009D (acknowledgments)
ACKNOWLEDGEMENTS We acknowledge support from the National Science Foundation through grant NSF-1108957. Support for Program numbers 12634 and 12575 was provided by NASA through a grant from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contract NAS5-26555. AIZ acknowledges support from NSF grant AST-0908280 and NASA grant ADP-NNX10AD47G. TC and MD acknowledge partial support of NASA grant HST- GO-12065.01-A. The archival and Guest Observer data are based on observations made with the NASA/ESA Hubble Space Telescope, which is operated by the Space Telescope Science Institute.
Paper: 2016A&A...587A..57Z (fulltext)
are compatible with the literature data and achieve an accuracy similar or better with respect to the published measurements, which is a good test for this new instrument. The orbital fit is a minimum fitting, as described in Esposito et al. (2013) , where all the four orbits are simultaneously taken into consideration, using the assumption that the planet periods are commensurate. The IFS measurement is used for the closest planet e, as it provides smaller error bars. We assumed circular and non-coplanar orbits and tested three different configurations for the orbital period ratios between the planets. The fitted orbits are shown in Fig. 5 . The parameters of the orbital solutions are summarized in Table 7 . The uncertainties on the parameters were estimated using a quadratic combination of statistical and systematic errors. The statistical errors stem from the sensitivity of the fit to the errors bars on the astrometric measurements and are derived from 1000 random trials of the astrometric measurements assuming normal distributions ( Maire et al. 2015 ). The systematic errors are related to the uncertainties on the stellar mass and distance (Table 7 ), which propagate into the derivation of the semimajor axis. To assess them, we performed two sets of orbital fits. The first set assumes the nominal values for the stellar mass and distance, and the second set assumes a combination of values at the edge of the derived ranges, which provides the largest offset in orbital periods. Two orbital solutions are favored with similar reduced , with the orbital period ratios 2d:1e and 3d:1e. The IFS data point for planet e is out of the orbital solution with the orbital period ratio 5d:2e at 2.5 significance level. Further astrometric monitoring is required to distinguish between the orbital period ratios 2d:1e and 3d:1e. Fig. 5 Relative astrometry for the planets around HR 8799. The points available from the literature are represented with black filled circles, while new SPHERE measurements are represented with purple filled squares. The solid blue lines represent the orbital solution labeled A in Table 7 . Red lines connect the predicted and observed positions for all the data points. The dotted and dashed lines in the panel for HR 8799 e indicate the orbital solutions labeled B and C in Table 5, respectively. Table 7 Orbital elements fitted on the astrometric data in Fig. 5 . Notes. Case A: orbital solution setting circular orbits and 8b:4c:2d:1e mean motion resonance. Case B: orbital solution setting circular orbits and 4b:2c:1d and 5d:2e mean motion resonances. Case C: orbital solution setting circular orbits and 4b:2c:1d and 3d:2e mean motion resonances. For each planet, the notations refer to the orbital period, inclination, longitude of the ascending node, eccentricity, argument of periapsis, time of periapsis passage, and semimajor axis. The error bars on the orbital parameters are at 1 (see text). We also indicate the reduced for the nominal