adsabs/WIESP2022-NER huggingface dataset:
Paper: 2017MNRAS.470.4925G (fulltext)
Table 6 shows a comparison of the jet half-opening angles derived from the best fits in this work with hydrodynamical modelling performed by Ryan et al. ( 2015 ) for four GRBs common to both studies. Our model produces systematically narrower jets (most likely caused by the model attempts to fit the early-time luminosity) that are partially consistent with Ryan et al. ( 2015 ) in errors (e.g. GRBs 051016B and 060614), and where they are not (e.g. GRB 061006), they are broadly consistent to ∼2σ–2.5σ. Table 6. Table showing the half-opening angles (in radians) for four GRBs, calculated from f B = 1 − cos (θ j ). θ j values are from the global best fits of this work (uncertainties are 95 per cent confidence interval); θ 0 values are from Ryan et al. ( 2015 ). Comparing our results with that of Gompertz et al. ( 2014 ), we can see the inclusion of fallback accretion within the propeller model allows for an improvement in fitting the ‘tail’ of the fading afterglow. This can be seen in GRBs 051227, 060614 and 061006, where Gompertz et al. ( 2014 ) did not produce such good fits to the tail. Hence, fallback accretion is a necessary addition to the propeller model in order to fully explain the energetics and morphologies of SGRBEEs. Additionally, the extended model handles variability and flares within the data much more naturally than Gompertz et al. ( 2014 ) and copes with the early-time luminosity detected by BAT. The parameters derived from the fits in Fig. 6 are presented in Table 7 . We find that the magnetic fields derived from the fits are in the moderate to high end of the parameter space and that the sample generally has slow initial spins. The slow initial spins are most likely due to the additional fallback spinning the magnetar up and, therefore, the constraints on high initial spin rates is relaxed. This has an impact on the value of the magnetic field derived as the fit moves along the correlation between B and P i discussed in Gompertz et al. ( 2014 ). The sample fits also tend to favour massive discs and narrow jet opening angles. This is most likely due to the model extracting as much of the available energy as possible to fit the high luminosities at early times in the light curve, data which was not included in the fits of Gompertz et al. ( 2014 ). The values of ε, δ, η dip and η prop are widely distributed throughout the parameter space. The derived parameters are consistent with predictions for a magnetar (Giacomazzo Perna 2013 ; Mereghetti, Pons Melatos 2015 ; Rea et al. 2015 ) and are also consistent with the results in Gompertz et al. ( 2014 ). Table 7. Parameters derived from the best global fits to the SGRBEE sample (bold values in Table 5 ). Reported errors are 95 per cent. Values marked with an [L] have reached a parameter limit; those marked with [F] were fixed during fitting. The values are also presented to indicate goodness of fit. We will now examine how increasing the number of free parameters affects the fits
Paper: 2018MNRAS.479.3634M (acknowledgments)
ACKNOWLEDGEMENTS The authors are grateful to the anonymous referee for a thoughtful report that helped improve the clarity of this work. AM was partly supported by an ‘ERASMUS+ for traineeship’ grant and partly by the NewCompStar network, COST Action MP1304. AM thanks the hospitality of the Anton Pannekoek Institute for Astronomy in Amsterdam, where most of this research was carried out. ND is supported by a Vidi grant awarded by the Netherlands Organization for Scientific Research (NWO). RW is supported by Top-1 grant from NWO. The authors acknowledge financial contribution from the agreement ASI-INAF I/037/12/0, as well as fruitful discussion with the international team ‘The disc–magnetosphere interaction around transitional millisecond pulsars’ at the International Space Science Institute, Bern. This work is partially supported by the HERMES Project, financed by the Italian Space Agency(ASI) Agreement n. 2016/13 U.O
Paper: 2016MNRAS.457.2301S (acknowledgments)
This work was supported by the National Science Foundation under grant no. AST-1313169. RM was also supported by an Alfred P. Sloan Fellowship. We thank Jonathan Blazek, Uroš Seljak, Robert Lupton, Elisa Chisari, and Shadab Alam for useful discussions about this work. We thank the anonymous referee for helpful suggestions, and we thank Ananth Tenneti for sharing the MB-II shape catalogue. Funding for SDSS-III has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, and the U. S. 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: 2017ApJ...834..185K (acknowledgments)
K.E.K. and M.P.E. were supported in part by NASA via the Air Force Research Laboratory. Support for G.C.S. was provided by NASA through contract No. 1257184 issued by the Jet Propulsion Laboratory (JPL), California Institute of Technology under NASA contract 1407. O.C.J. acknowledges support from NASA grant NNX14AN06G. This research has made use of NASA’s Astrophysical Data System Bibliographic Services, the Simbad and VizieR databases operated at the Centre de Données astronomiques de Strasbourg, and the Infrared Science Archive at the Infrared Processing and Analysis Center, which is operated by JPL. We thank the JWST Helpdesk and the MIRI team for making the MIRI filter functions available and the referee for helpful suggestions to clarify the paper.
Paper: 2018ApJ...867..129C (acknowledgments)
We acknowledge Venus Express and Pioneer Venus Orbiter teams as well as the NASA Planetary Data System ( http://ppi.pds.nasa.gov ) for providing the data. We also thank ACE team and NASA CDAWeb for providing magnetometer data at 1 au, and the World Data Center SILSO ( http://www.sidc.be/silso/ ) for using the sunspot number data. This work is supported by the Science and Technology Development Fund of Macao SAR (008/2016/A1 and 119/2017/A3) and National Natural Science Foundation of China (NSFC) under grants 41564007 and 41731067.
Paper: 2020MNRAS.497.5118S (fulltext)
to a time-span of 3.7 yr). The magnitude difference is only 0.08 mag, which is within the combined photometric error (0.11 mag). We also find that the I -band magnitude at t = 586 d is consistent with the extrapolation of SN 2014C’s earlier MMT light curve. Thus, it is possible that SN 2014C has faded significantly in the I band by t = 586 d, after which the light from Cluster A dominates the I -band brightness. We also note that the emission lines in the I band are much weaker than those in the B, V , and R bands ( M15 ; Anderson et al. 2017 ). To be conservative, however, we do not exclude the possibility that the SN was still bright in the I band and we regard the measured F814W brightness as an upper limit for Cluster A . Similarly, the F336W -band magnitude also remains almost constant from t = 1003 to 1462 d (which corresponds to a time-span of 1.3 yr). The magnitude difference is only 0.06 mag, which is comparable to the combined photometric error (0.05 mag). Note that during the same period, the F555W -band magnitudes faded by 0.17 mag. Yet, SN 2014C, like many other interacting SNe, exhibits a blue pseudo-continuum, which may arise from blended iron lines produced by the ejecta–CSM interaction ( M15 ; see also, e.g. SN 2006jc, Pastorello et al. 2007 ; Bufano et al. 2009 ). As a result, we cannot rule out the possibility that the F336W band is contaminated by the pseudo-continuum emission from the SN. Thus, we regard the F336W -band brightness also as an upper limit for Cluster A . The F275W and F300X bands were only observed at a single epoch. Both bands cover the [Mg ii ] doublets at ∼2800 Å; the F300X band also covers the [Fe ii ] line at ∼2950 Å. These two lines are important coolants and often appear to be very strong in emission for interacting SNe (Fesen et al. 1999 ; Fransson et al. 2005 ; Bufano et al. 2009 ). Thus, F275W and F300X bands may be contaminated by the SN light, although it is difficult to estimate the amount of its contribution. Still, they place important upper limits for the brightness of Cluster A at UV wavelengths. In summary, SN 2014C’s light was still significant until at least 4 yr post-explosion (the last epoch of F555W observation). This is consistent with the conclusion of T19 that SN 2014C experiences ongoing interaction between its ejecta and CSM at large distances. Still, the HST observations provide very important constraints on Cluster A ’s SED from the UV to the F814W band, a much wider wavelength coverage than the pre-explosion observations in the B, V, R , and F658N bands. Thanks to the higher spatial resolution provided by the new data, the measured SED of Cluster A is now much cleaner, i.e. without significant contamination from its nearby sources. 4 PROPERTIES OF THE HOST STAR CLUSTER We further try to infer the properties of Cluster A by fitting model spectra to its observed SED. Note that Cluster A is still associated with a considerable amount of (ionized) gas, given its bright
Paper: 2019MNRAS.489.4632R (fulltext)
of the solution operator is Note that the exponential in equation ( 6 ) is simply a formal way to write down an operator. It is acting on the function to its right, in the above case this is the identity Id. To see that this construct indeed provides a solution to the differential equation with initial conditions y 0 , note that the series expansion of the exponential results in a series of differential operators acting on Id( y 0 ). This series is simply the Taylor expansion of the solution φ around t = 0. Since it is not possible to write down an explicit solution for the N -body problem, we turn to operator splitting methods. One way to view an operator splitting method is to take a differential equation of the form and rewrite it as Note that due to the linearity of the Poisson bracket, we have . This implies that for Hamiltonian systems we can split the equations of motion by splitting the Hamiltonian H into two parts, i.e. H = A + B . Depending on the splitting, we get two new sets of differential equations that we can study independently: In practice, we will choose a splitting that allows us to easily solve these new differential equations independently. In other words, we want to be able to explicitly write down the solution operators and . Using the notation from above, we have but because these exponentials are operators, not scalars, we unfortunately cannot use standard rules for exponentials. Specifically, we typically have For illustration purposes, imagine the above equation were true. We could then solve the differential equation in equation ( 7 ) by first applying the solution operator φ [ B ] to our initial conditions for a time t (remember, operators act to the right), then take the result and apply the solution operator φ [ A ] on it, again for time t . The result of the last operation would then correspond to the solution φ [ H ] at time t . But because the equation above is not true, we need to rely on the Baker–Campbell–Hausdorff (BCH) identity, 2 to get a similar expression where is a commutator. So if we apply the splitting procedure described above (corresponding to the left-hand side of equation 12 ), we do not solve the differential equation , but some other differential equation (given by the argument in the exponential on the right-hand side of equation 12 ). Formally, we can write down the Hamiltonian corresponding to the differential equation that we are actually solving as where we have used which follows from the Jacobi identity. Although we do not end up solving the Hamiltonian system we were trying to solve, the fact that we can write down a different Hamiltonian system that we seem to solve exactly is often quoted as why symplectic integrators have superior properties. 3 We can study the properties of this ‘nearby’ Hamiltonian to see how far our approximation is from the true solution. As we can see from the equation above, for small t , the Hamiltonian is close to the original Hamiltonian and only
Paper: 2017A&A...607A..56R (acknowledgments)
We thank the anonymous referee for the constructive report that helped us to improve the readability of the paper. The author acknowledges the use of the TOPCAT 3 and STILTS 4 software packages written by Mark Taylor (2005, 2006). This paper uses data taken with the MODS spectrographs built with funding from NFS grant AST-9987045 and the NSF Telescope System Instrumentation Program (TSIP), with additional funds from the Ohio Board of Regents and the Ohio State University Office of Research. We acknowledge financial support from the ARCHES project (7th Framework of the European Union, No. 313146). Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the US 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 on observations obtained with XMM-Newton , an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA. This research has made use of data obtained from the 3XMM XMM-Newton serendipitous source catalogue compiled by the 10 institutes of the XMM-Newton Survey Science Centre selected by ESA.
Paper: 2016A&A...589A.126V (acknowledgments)
We thank the anonymous referee for the comments that have greatly helped improve the presentation of our results. We also thank Emanuele Dalessandro for sending us his HST photometry for M3, and Luca Sbordone and Fiorella Castelli for their expert advice in the use of the ATLAS9 and SYNTHE packages. Support for A.A.R.V., M.C., J.A.-G., R.C.R., and S.A. is provided by the Ministry for the Economy, Development, and Tourism’s Programa Iniciativa Científica Milenio through grant IC 120009, awarded to the Millennium Institute of Astrophysics (MAS); by Proyecto Basal PFB-06/2007; and by Proyecto FONDECYT Regular #1141141. A.A.R.V., J.A.-G., and R.C.R. acknowledge support through FONDECYT postdoctoral grants #3140575, 3130552, and 3130320, respectively. J.A.-G. acknowledge support by the FIC-R Fund, allocated to the project 30321072. This work was done (in part) using the Geryon computer at the Center for Astro-Engineering UC, part of the Basal PFB-06 project, which received additional funding from QUIMAL 130008 and Fondequip AIC-57 grants for upgrades. Additional support for this project is provided by CONICYT’s PCI program through grant DPI20140066.
Paper: 2017ApJ...842...65Z (acknowledgments)
We thank the anonymous referee for the useful comments. We also thank Terry Mahoney for revising the English of this paper. This work is based on observations made with the Gran Telescopio Canarias (GTC), installed at the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofísica de Canarias, on the island of La Palma. This paper is also based on observations made with ESO Telescopes at the Paranal Observatory under program ID 090.C-0766 and also includes data gathered with the 6.5 m Magellan Telescopes at Las Campanas Observatory, Chile. This research has benefited from the SpeX Prism Spectral Libraries 11 11 http://pono.ucsd.edu/adam/browndwarfs/spexprism and the Montreal Brown Dwarf and Exoplanet Spectral Library, 12 12 https://jgagneastro.wordpress.com/the-montreal-spectral-library/ maintained by Adam Burgasser and by Jonathan Gagné, respectively. This research is partly financed by the Spanish Ministry of Economy and Competitivity through projects AYA2014-54348-C3-2-R, AYA2015-69350-C3-2-P, AYA2016-79425-C3-2-P, and the Chilean FONDECYT Postdoctoral grant 3140351. K.P.R. acknowledges CONICYT PAI Concurso Nacional Inserción en la Academia, Convocatoria 2016 Folio PAI79160052. Facilities: VLT (ISAAC), GTC (OSIRIS), Magellan (FIRE).
Paper: 2017A&A...608A..17T (acknowledgments)
We are grateful to Ilia Potravnov for a number of useful comments and to the Swift team for the execution of our ToO request. This work was supported by the Russian Science Foundation grant 14-12-01287 (S.S.T., A.A.M., A.A.L., V.S.), the Academy of Finland grant 268740, and the National Science Foundation grant PHY-1125915 (J.P.). V.D. thanks the Deutsches Zentrums for Luft- und Raumfahrt (DLR) and the DFG for financial support (grant DLR 50 OR 0702). We also acknowledge the support of COST Action MP1304.
Paper: 2018MNRAS.478.2851M (acknowledgments)
ACKNOWLEDGEMENTS Our anonymous referee provided valuable feedback that helped improve this manuscript. We thank Hiroyuki Hirashita for providing us with the tabulated form of grain velocities calculated over an extended grain-size range and used in Hirashita Yan ( 2009 ). We also thank Takaya Nozawa for making available the tabulated versions of the SN grain-size distributions presented in Nozawa et al. ( 2007 ) and the SN dust destruction efficiencies used in Asano et al. ( 2013b ). Finally, we are grateful to Volker Springel for sharing access to arepo . The simulations were performed on the joint MIT-Harvard computing cluster supported by MKI and FAS. RM acknowledges support from the DOE CSGF under grant number DE-FG02-97ER25308. MV acknowledges support through an MIT RSC award and the support of the Alfred P. Sloan Foundation. PT acknowledges support from NASA through Hubble Fellowship grant HST -HF2-51341.001-A awarded by STScI, which is operated under contract NAS5-26555. RK acknowledges support from NASA through Einstein Postdoctoral Fellowship grant number PF7-180163 awarded by the Chandra X-ray Center, which is operated by the Smithsonian Astrophysical Observatory for NASA under contract NAS8-03060.
Paper: 2019MNRAS.487.2742D (acknowledgments)
ACKNOWLEDGEMENTS We thank the referee for his/her constructive reports and very helpful suggestions regarding the presentation of this paper and the interpretation of our results, S. J. Aarseth for providing the code used in this research, O. Vaduvescu, F. Roig, M. N. De Prá, and S. Deen for comments, and A. I. Gómez de Castro for providing access to computing facilities. This work was partially supported by the Spanish ‘Ministerio de Economía y Competitividad’ (MINECO) under grants ESP2015-68908-R and ESP2017-87813-R. In preparation of this paper, we made use of the NASA Astrophysics Data System and the Minor Planets Center (MPC) data server.
Paper: 2019MNRAS.487.3270P (fulltext)
for the Galactic bar (Dehnen 2000 ). For all of the following calculations, we assumed a rotation velocity of 220 km s −1 (Kerr Lynden-Bell 1986 ; Bovy et al. 2012 ), a distance to the Galactic centre of 8.2 kpc (Bland-Hawthorn Gerhard 2016 ), and the same peculiar solar motion as in Section 3 , and we integrated forward in time up to 1 Gyr. The galpy software allows one to calculate parameters of individual orbits, including e.g. eccentricity e , maximum height above or below the Galactic plane z max , and peri and apocenter of the orbit, r peri and r apo , respectively. For each star, we ran a Monte Carlo simulation. We calculated 200 orbital solutions in which we varied the initial conditions (proper motions, radial velocities, and distances) within their uncertainties assuming that their distribution is Gaussian. From the resulting distribution, we calculated the median orbital parameters for both Oosterhoff groups. The results are listed in Table 2 . In order to test the differences between the two Oosterhoff groups we used the CM test (see Section 3 ). For the entire distribution of orbital parameters z max , e, r peri , and r apo , the CM test yields p CM = 0.68, there is no statistical evidence that the two distributions are not identical. Table 2. Median values of calculated parameter orbits for the individual Oosterhoff groups. Column 1 lists the studied components, columns 2 and 3 list values for each Oosterhoff group. Fig. 5 depicts the distribution of apocentric distances for the studied RR Lyrae stars. In this diagram, we see that both Oosterhoff groups have very similar apocentric distance distributions. We also noticed a sudden decrease at ≈ 3 kpc (log( r apo ≈ 0.5 kpc), which is roughly in agreement with the overlap between bulge, metal-weak thick disc, and inner halo (3.5 kpc; Ness et al. 2013b ). In fact, the majority of the RR Lyrae stars in our sample stay inside the bulge volume during their full orbits. Figure 5. The distribution of apocentric distances for the studied variables. The blue and red histograms represent the Oo I and Oo II groups, respectively. Based on the negligible difference in the median value of z max between the two Oosterhoff groups, we explored whether the Oo II variables could follow higher orbits from the Galactic plane and thus possibly come from the halo. Based on the distribution of apocentric distances (Fig. 5 ), as a nominal boundary for the Galactic bulge we chose z max > 3 kpc (roughly above ± 20 deg at the distance of the Galactic bulge) and calculated the number of stars that fulfil this condition in at least 50 per cent of the iterations of the Monte Carlo simulation. In total, only 34 RR Lyrae stars fulfil the aforementioned condition. Among these 34 variables, 35 per cent belong to the Oo II population. This is a significant increase in comparison with the overall bulge and halo Oosterhoff population where the fraction of Oo II pulsators is around 25 per cent (Drake et al.