adsabs/WIESP2022-NER huggingface dataset:
Paper: 2019MNRAS.490.1652G (fulltext)
positions, which are greatly improved by refining PECs using experimentally measured line positions, this is seldom true for the line intensities, for which DM and TDM curves are usually best computed ab initio (Tennyson 2014 ). The line intensities, which are directly based on the ab initio DMCs, often suffer from the numerical noise, especially those from the overtone transitions (Medvedev et al. 2016 ). In order to reduce this noise, it was sufficient to apply a dipole moment cut-off of 10 −8 D to the vibrational transition matrix elements, see e.g. Yurchenko et al. ( 2018c ). When computing these matrix elements, the original ab initio X – X and A – X (T)DMCs were mapped on the duo grid using the standard cubic spline interpolation technique. The final spectroscopic model in the form of the duo input file is provided as part of the supplementary data to this paper and can be also found at www.exomol.com . This also includes our fitting set of experimental frequencies. A sample of the refinement to the data of Ramsay ( 1952 ) for the X 2 Π – A 2 Σ + system is shown in Table 2 . The Obs. − Calc. values of the fit are comparable to the actual experimental measurement uncertainty. Although these may seem from the outset to be of much lower accuracy compared to previous ExoMol line lists for non-transition metal-containing diatomics, given that in the regime of ≈30 000 cm −1 an uncertainty of 0.5 cm −1 corresponds to a resolution R of ≈60 000, which is within the realms of high-spectral resolution methods of characterizing exoplanet atmospheres (Snellen 2004 ; de Kok et al. 2014 ). For all isotopologues of SH-type ( 33 SH, 34 SH, 36 SH), we use the same empirical model (PECs, SOCs, EAMCs, SRCs, and BOBCs) developed for 32 SH. When solving the Schrödinger equation, the mass of sulphur is replaced by the nuclear mass of the corresponding isotope. For the isotopologue 32 SD, however, this approach usually leads to too large discrepancies with the experimental line positions. We have therefore refined the 32 SH model by fitting to the experimental line positions of 32 SD from the A – X band, which were taken from the same sources listed in Table 1 . Ideally, this refinement should be only applied to the BOB-correction. However, due to the lack of the experimental data on the X 2 Π state for SD, we could not build an accurate and self-consistent model without including other curves into the fit. We have therefore varied the parameters of PEC( A ), PEC ( B ), BOBC ( X ), BOBC ( A ), SRC ( X ), and SRC ( A ) by fitting to the SD line positions from Ramsay ( 1952 ), Johns Ramsay ( 1961 ), and Pathak Palmer ( 1969 ). All other curves were fixed to the those from the spectroscopic model of 32 SH, including PEC( X ), SOC ( X ), SOC( X – A ), SOC( X – B ), SOC( A – B ), EAMC( X – A ), and EAMC( A – B ). The rms error of this fit to the A – X line positions if SD is 0.4 cm −1 . The spectroscopic model for SD is included in the supplementary material as a
Paper: 2015MNRAS.448.1644S (acknowledgments)
The authors would like to thank Daniel Gifford and Doron Lemze for helpfully answering questions regarding the caustic technique and useful comments. We also thank the anonymous referee for constructive comments which helped to improve the work. The Dark Cosmology Centre is funded by the Danish National Research Foundation. RW acknowledges support through the Porat Postdoctoral Fellowship.
Paper: 2016ApJ...824L..22M (fulltext)
in this limit. is given bywhere n , a , e , i are a resonant eKBO’s mean motion, semimajor axis, eccentricity, and inclination (relative to HP’s orbit plane), is the planet’s semi-minor axis, and c is a numerical coefficient whose asymptotic value is ∼5.8 for a ≪ a ′ ( c will be larger for a approaching a ′). For numerical estimates we adopt e ′ = 0.2 and c = 8. We find that for m ′ = 10 M ⊕ , the small amplitude oscillation periods for the four eKBOs are 0.6–0.94 Gyr, a factor of a few smaller than the age of the solar system. We conclude that 10 M ⊕ is approximately the minimum mass necessary to maintain periodic orbits of the third kind for the four eKBOs. 2.3. Planet’s Current Location in Its Orbit Each resonant eKBO constrains the longitude of HP because the critical resonant angle must librate about either 0° or 180°. For illustration, consider the libration of about 180° for Sedna’s 3/2 resonant orbit. Given that Sedna is presently near its perihelion, λ ≈ π , the planet can be in proximity to one of three longitudes: 60°, 180°, or 300° from Sedna’s longitude of perihelion, π . The trace of the resonant orbit in the rotating frame (Figure 2 ) shows that libration of φ would allow the planet to be located as far as ∼54° from these longitudes. This means that only three small zones of ±6° around each of three locations, π , and are excluded for the planet’s present location. For e ′ = 0.2, we find that the excluded zone is a little larger, ±75 around each of the same three locations. However, such geometric estimates of exclusion zones do not account for the planet’s gravity. To account for this, we computed surfaces of section for the circular planar restricted three body problem for several values of the planet mass in the range of 10 −6 –10 −3 M ⊙ , for the neighborhood of the 3/2, 5/2, 3/1, and 4/1 MMRs and the extreme eccentricities of the four eKBOs (0.85, 0.87, 0.85 and 0.7, respectively). Figure 4 shows a sample of these for planet mass 3 × 10 −5 M ⊙ (approximately 10 M ⊕ ). From these surfaces of section, we determine the range of excluded longitudes of the planet by measuring the complement of the large libration zones of the 3/2 and 5/2 MMRs. We also conservatively exclude a small range, ±5°, around the homoclinic points of the 3/1 and 4/1 resonant eKBOs. Figure 5 shows our estimates of the excluded zones, projected on the sky, for a planet of mass 3 × 10 −5 M ⊙ ; we plot a representative case of the planet’s low inclination orbit (in green) and also a representative case of its high inclination orbit (in red). (Note that a range of orbital planes near each of these representative cases is allowed by our calculations.) We find that approximately half of the ecliptic longitudes are excluded in both cases; the excluded zone is not contiguous. Figure 4. Surfaces of section for the circular planar restricted three body problem, for mass ratio , and Jacobi constant values as indicated. These are phase space portraits in the
Paper: 2019ApJ...884...53F (fulltext)
is of similar order to earlier UV-based measurements (∼1 M ⊙ yr −1 ; Shull et al. 2009 ; Lehner Howk 2011 ; R17 ), but higher than radio-based measurements derived from large 21 cm HVCs (≈0.08 M ⊙ yr −1 ; Putman et al. 2012 ) and higher than the flow rates derived for individual clouds like Complex C (≈0.1 M ⊙ yr −1 ; Wakker et al. 1999 ; Collins et al. 2003 ; Shull et al. 2011 ). This makes sense because most HVC mass is in ionized HVCs (Lehner Howk 2011 ). The ratio of the UV inflow rate to the 21 cm inflow rate quantifies the relative mass in ionized and neutral gas. Our inflow rate calculations indicate (H ii )/ N (H i ), equivalent to a mean HVC ionization fraction f H II = N (H ii )/ N (H i +H ii ) = 0.87. This is higher than the mean ionization fraction reported in large 21 cm HVCs such as Complex C and Complex A because such large clouds are atypical; most HVCs are ionized and not detected at 21 cm (Shull et al. 2009 ; Lehner Howk 2011 ; R17 ). For comparison with the HVC mass flow rates, the intermediate-velocity mass flow rates are of interest. These can be derived from optical observations of Ca ii and Na i absorption (Ben Behkti et al. 2012 ). The Ca ii and Na i doublets are much less saturated than UV resonance lines, particularly at intermediate velocities, resulting in cleaner measurements. Ben Behkti et al. ( 2012 ) report a population of Ca ii IVCs with a sky covering factor of ≈30%, a mean v GSR = −42 km s −1 , and an associated N (H i ) distribution peaking near log N (H i ) ≈ 19. Combining these numbers with a typical IVC distance of 2 kpc (Richter et al. 2017 ) we derive an H i mass flow rate of ≈0.04 M ⊙ yr −1 in these Ca ii IVCs. If a similar ionization correction N (H ii )/ N (H i ) = 6.6 applies to IVCs as to HVCs, then IVCs represent a total mass flow rate of ≈0.27 M ⊙ yr −1 , comparable to the HVC flow rate (this similarity occurs because the larger column densities of IVCs are compensated by their smaller velocities and distances). To fully explore this, a more thorough investigation of the IVC mass flow rates using UV data would be worthwhile, including a distinction into inflow and outflow and detailed ionization modeling. We emphasize that the HVC mass flow rates derived in this paper are effectively instantaneous , as they represent the mass flux contained in HVCs observed at the current epoch. Hydrodynamic simulations (Heitsch Putman 2009 ; Armillotta et al. 2017 ) show that HVC lifetimes are typically a few 100 Myr depending on HVC mass and the density contrast with the surrounding hot medium, after which the clouds will become disrupted and evaporated by coronal interactions. Caution should therefore be used in integrating our mass flow rates over long timescales to calculate the total mass delivered in and out of the disk, as the gas will change phase over time via evaporation and subsequent condensation (e.g., Fraternali 2017 ; Voit et al. 2019 ). Finally, we can use our HVC mass outflow rate measurement to
Paper: 2019MNRAS.484.5049M (acknowledgments)
ACKNOWLEDGEMENTS FEM and EEM acknowledge support from National Science Foundation (NSF) award AST-1313029. Part of this research was carried out at the Jet Propulsion Laboratory, California Institute of Technology, under a contract with the National Aeronautics and Space Administration (NASA). EEM acknowledges support from the Jet Propulsion Laboratory Exoplanetary Science Initiative and the NASA NExSS Program. KL was supported by grant AST-1208239 from the NSF. We thank the staff of CTIO and the Blanco 4-m telescope in particular for their help and hospitality during the 2013 and 2015 observing runs. We thank Frank Valdes and Robert Gruendl for their help understanding the DECam community pipeline and calibration of photometry and astrometry. The Center for Exoplanets and Habitable Worlds is supported by the Pennsylvania State University, the Eberly College of Science, and the Pennsylvania Space Grant Consortium. We thank the anonymous referee for a very thorough report which significantly improved the paper.
Paper: 2016A&A...595A..28C (acknowledgments)
The authors would like to thank Alan Calder for many fruitful discussions and exchanges. Part of the software used in this work was developed by the DOE-supported ASC/Alliances Center for Astrophysical Thermonuclear Flashes at the University of Chicago. This work has been partially supported by the Spanish MINECO grant AYA2014–59084–P, by the E.U. FEDER funds, by the AGAUR/Generalitat de Catalunya grant SGR0038/2014, and by the U.S. Department of Energy, Office of Science, Office of Nuclear Physics. We also acknowledge the Barcelona Supercomputing Center for a generous allocation of time at the MareNostrum supercomputer.
Paper: 2019MNRAS.489.1837R (acknowledgments)
ACKNOWLEDGEMENTS We thank an anonymous referee for stimulating comments and Elisabeta Lusso and Gabriele Ghisellini for sharing results of previous publications. The data collected by the WEBT collaboration are stored in the WEBT archive at the Osservatorio Astrofisico di Torino – INAF ( http://www.oato.inaf.it/blazars/webt/ ); for questions regarding their availability, please contact the WEBT President Massimo Villata ( massimo.villata@inaf.it ). This publication uses data obtained at Metsähovi Radio Observatory, operated by Aalto University in Finland. We acknowledge financial contribution from the agreement ASI-INAF n.2017-14-H.0 and from the contract PRIN-SKA-CTA-INAF 2016. PR and SV acknowledge contract ASI-INAF I/004/11/0. We acknowledge support by Bulgarian National Science Programme ‘Young Scientists and Postdoctoral Students 2019’, Bulgarian National Science Fund under grant DN18-10/2017 and National RI Roadmap Projects DO1-157/28.08.2018 and DO1-153/28.08.2018 of the Ministry of Education and Science of the Republic of Bulgaria. GD and OV gratefully acknowledge the observing grant support from the Institute of Astronomy and Rozhen National Astronomical Observatory, Bulgarian Academy of Sciences via bilateral joint research project ‘Study of ICRF radio-sources and fast variable astronomical objects’ (head – G.Damljanovic). This work is a part of the Projects No. 176011 (‘Dynamics and Kinematics of Celestial Bodies and Systems’), No. 176004 (‘Stellar Physics’), and No. 176021 (‘Visible and Invisible Matter in Nearby Galaxies: Theory and Observations’) supported by the Ministry of Education, Science and Technological Development of the Republic of Serbia. This research was partially supported by the Bulgarian National Science Fund of the Ministry of Education and Science under grants DN 08-1/2016, DN 18-13/2017, and KP-06-H28/3 (2018). The Skinakas Observatory is a collaborative project of the University of Crete, the Foundation for Research and Technology – Hellas, and the Max-Planck-Institut für Extraterrestrische Physik. The St Petersburg University team acknowledges support from Russian Science Foundation grant no. 17-12-01029. The Abastumani team acknowledges financial support by the Shota Rustaveli National Science Foundation under contract FR/217950/16. This work was partly supported by the National Science Fund of the Ministry of Education and Science of Bulgaria under grant DN 08-20/2016, and by funds of the project RD-08-37/2019 of the University of Shumen. The Astronomical Observatory of the Autonomous Region of the Aosta Valley (OAVdA) is managed by the Fondazione Clément Fillietroz-ONLUS, which is supported by the Regional Government of the Aosta Valley, the Town Municipality of Nus and the ‘Unité des Communes valdôtaines Mont-Émilius’. The research at the OAVdA was partially funded by two ‘ Research and Education’ grants from Fondazione CRT. This research has made use of the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.
Paper: 2015MNRAS.450.1777K (acknowledgments)
We thank Y. Qian and P. Banerjee for help with the torch network code. DK is supported in part by a Department of Energy Office of Nuclear Physics Early Career Award, and by the Director, Office of Energy Research, Office of High Energy and Nuclear Physics, Divisions of Nuclear Physics, of the US Department of Energy under contract no. DE-AC02-05CH11231. RF was supported by NSF Division of Astronomical Sciences collaborative research grant AST-1206097, and a UC Office of the President grant. BDM gratefully acknowledges support from the NSF grant AST-1410950 and the Alfred P. Sloan Foundation. This work was supported in part by NSF grant no. PHYS-1066293 and the hospitality of the Aspen Center for Physics. The software used in this work was in part developed by the DOE NNSA-ASC OASCR Flash Center at the University of Chicago. We are grateful for computing time made available by the National Energy Research Scientific Computing Center, which is supported by the Office of Science of the US Department of Energy under contract no. DE-AC02-05CH11231. Computations were performed using Carver and Hopper .
Paper: 2018MNRAS.475..999B (fulltext)
the behaviour of emission lines at 0.913 keV (Ne ix ) and 1.02 keV (Ne x ) with the pulse phase. We summarize the results as follows: Light curve obtained using the EPIC -pn data during its current spin-up phase showed dips. Unlike other flaring sources like LMC X–4 and SMC X−1, it is interesting to notice a broad dip in the light curve soon after the decay of a large flare. This feature is similar to that observed in the bursting pulsar, GRO J1744−28. The light curve of GRO J1744−28 showed the presence of a dip and recovery period following each outburst (see e.g. Giles et al. 1996 ) and the X-ray spectrum of GRO J1744−28 showed no significant change going from quiescence to outburst (Cannizzo 1996 ). The same authors proposed that the outbursts observed in the bursting pulsar could be due to Lightman–Eardley instability (Lightman Eardley 1974 ) in the accretion disc and the material that is evacuated on to the pulsar during an accretion event is replenished by material flowing in from further out; hence, the dip and recovery in the light curve following an outburst. After performing intensity-resolved spectroscopy of 4U 1626−67, we found that overall there is no change in the shape of shape of the spectrum (see Fig. 9 ) except that the continuum and line parameters follow an increasing trend with intensity. Therefore, it is plausible that a similar mechanism might be responsible for the presence of flares and sharp dips in the light curve of 4U 1626−67. A QPO feature around 3 mHz is observed in the PDS. This feature has been observed for the first time using X-ray data of current spin-up phase. The feature at 3 mHz also shows strong energy dependence. The feature is sharp at low energies as flares dominate at low energies. The energy dependence of fractional rms amplitude of a QPO has been used as a tool to understand the physical origin of the QPO (see e.g. Gilfanov, Revnivtsev Molkov 2003 ; Cabanac et al. 2010 ; Mukherjee Bhattacharyya 2012 ). The fractional rms amplitude of a 3 mHz QPO observed in 4U 1626−67 showed an increase up to 5 keV and thereafter its value saturates (probably due to lower count rates at higher energies). Therefore, it seems that fluctuations in the blackbody component could be a plausible cause of the observed mHz QPO in 4U 1626−67. We found that the pulse profile shape below 2 keV is different from that seen during spin-down phase of 4U 1626−67 (Fig. 4 ). Moreover, during its spin-up phase, pulse profiles below 2 keV are quite different from that seen above 2 keV (Fig. 3 ). A possible explanation to these observations is changes in the emission diagram of the accretion column. During the low-luminosity phase (spin-down) of 4U 1626−67, the emission of the accretion column is concentrated in a beam, oriented along the magnetic field axis while during the high-luminosity phase (spin-up) the emission diagram changed to the fan beam pattern (Basko Sunyaev 1975 ). Soft X-ray emission (below 2 keV) is attributed to
Paper: 2019A&A...632A...3G (acknowledgments)
We thank Arash Bahramian, Nate Bastian, Robert Mathieu, Wolfram Kollatschny, Stan Lai, and Mark Gieles for helpful discussions. BG, SD, SK and PMW acknowledge support from the German Ministry for Education and Science (BMBF Verbundforschung) through grants 05A14MGA, 05A17MGA, 05A14BAC, and 05A17BAA. AA is supported by the Carl Tryggers Foundation for Scientific Research through the grant CTS 17:113. SK gratefully acknowledges funding from a European Research Council consolidator grant (ERC-CoG-646928- Multi-Pop). JB acknowledges support by FCT/MCTES through national funds by grant UID/FIS/04434/2019 and through Investigador FCT Contract No. IF/01654/2014/CP1215/CT0003. This research is supported by the German Research Foundation (DFG) with grants DR 281/35-1 and KA 4537/2-1. Based on observations made with ESO Telescopes at the La Silla Paranal Observatory under programme IDs 094.D-0142, 095.D-0629, 096.D-0175, 097.D-0295, 098.D-0148, 0100.D-0161, 0101.D-0268, 0102.D-0270, and 0103.D-0204. Based on observations made with the NASA/ESA Hubble Space Telescope, obtained from the data archive at the Space Telescope Science Institute. STScI is operated by the Association of Universities for Research in Astronomy, Inc. under NASA contract NAS 5-26555. Supporting data for this article is available at: http://musegc.uni-goettingen.de .
Paper: 2019A&A...628A..82P (fulltext)
). Stars and SPH particles were coupled by direct point-to-point interaction, without any approximation for the gravitational field. Finally, we corrected the gas positions and velocity according to Eq. ( A.4 ) by recalculating the accelerations and the energy rates at the new stage n + 1. Appendix B: Technical features B.1. Neighbor search and acceleration updating Before we employed it to compute gravitational interactions, the tree grid was also used to support the operations related to the nearest-neighbor search. To calculate the density and the smoothing length, each i th particle starts with a first-guess value of h i and makes a tree walk by adopting an opening criterion that is slightly different from Eqs. ( 20 ) and ( 21 ). Given a generic particle of index i, in order to find the other points enclosed in its SPH kernel support, we checked for each cube the overlap with a sphere of radius 2 h i centered onto the point r i = ( x i , y i , z i ). A box is opened only if the following three conditions are valid: where ( x A , y A , z A ) are the coordinates of the geometrical center of the cube and D L is its side length. The tree walk continues by opening the further cubes according to the last rule, until the single particles are reached and the neighborhood is thus determined. At the end of the walk, a temporary value of density ρ i is calculated. If the number of encountered neighbors differs from the expected number (and the difference exceeds the tolerance number), h i is updated according to Eq. ( 5 ). Then, a new tree walk is performed, and a new value for the density is computed. The tree walk and the h i , ρ i updating are executed cyclically until the number of neighbors converges on the desired value. During the tree walk, the various quantities P i , Ω i , ζ i , f , ∇ ⋅ v i , and ∇ × v i are computed at the same time together with the density. After this preliminary phase, the acceleration a and the thermal energy rate can be computed by performing a further tree walk for each particle. During the walk, the hydrodynamics contribution to the acceleration and the gravity field are computed at the same time. For each box, before Eq. ( 20 ) (or Eq. ( 21 )) is applied, the code primarily checks Eq. ( B.1 ) to distinguish cubes that might contain some neighbor particles for the SPH interpolation. If this is satisfied, the box is opened, otherwise the particles inside will not contribute at all to hydrodynamics. After this primary check, the algorithm performs a secondary control by applying the opening criterion of Eq. ( 20 ) (or Eq. ( 21 )) to decide whether the multipole approximation can be applied to the Newtonian field. As a result, practically, a generic i th particle will interact with its local SPH neighborhood following Eqs. ( 7 ) and ( 8 ), experiencing the Newtonian force through a direct particle-to-particle interaction. The remaining particles that lie outside the SPH domain contribute to the accelerations with or