adsabs/WIESP2022-NER huggingface dataset:
Paper: 2017MNRAS.465.4406K (fulltext)
BBHs in the sense that it can be detectable before the GW signal and that the luminosity increases with time. However, the luminosity and total energy are too low to explain GRBs or the GBM event (Connaughton etal. 2016 ). Figure 3. The photon fluxes from the internal shocks for models A, B and C. The 1-scan sensitivity of MAXI GSC is also plotted. It is found that all the models are detectable by MAXI GSC. The optical follow-up surveys for GW counterparts, such as the Panoramic Survey Telescope Rapid Response System (Pan-STARRS) and the Japanese collaboration for Gravitational wave ElectroMagnetic follow-up, have a sensitivity of 1921 mag (Morokuma etal. 2016 ; Smartt etal. 2016 ). Assuming that 10percent of L AG is in the optical range, we calculate the distance of detection limit for the afterglow, d L, lim , whose values are tabulated in Table 1 . While the afterglow is fainter than the internal shock emission, the optical follow-up surveys of afterglows can detect comparable or more distant events than the X-ray monitoring systems owing to the good sensitivity. However, d L, lim for model A is shorter than the distance of the observed GW events (Abbott etal. 2016a , b ). Although d L, lim is larger for IMBHs and MBHs, we do not discuss the detection probability because the merger rates of IMBH and MBH binaries are very uncertain. 4 SUMMARY and DISCUSSION We study evolution of an accretion disc in BBH systems and propose an evolutionary track of the disc, which leads to different conclusion from the previous work (Perna etal. 2016 ). At first, the disc viscously expands outward but the companion prevents the disc from expanding beyond r out due to the tidal torque. The evolution of viscous disc results in the decrease of the disc mass and the temperature. When the disc sufficiently cools down (typically less than 3000K), the dead disc forms because MRI becomes inactive. Since the thermal instability causes the rapid drop of the disc temperature, the disc becomes dead when the temperature becomes less than a few tens of thousands K. This dead disc remains until the binary separation sufficiently decreases. As the binary separation decreases, the position at which the tidal torque is effective moves inward, and the mass of the outer rim increases. Then, the angular momentum is transported by the tidal torque, which induces the mass inflow from the outer rim to the dead disc. When the mass inflow by the tidal torque becomes higher than , the accretion heating activates MRI, restarting the mass accretion from the disc to the central BH (the disc revival). This disc revival typically happens tens of thousands years before the merger event. The evolution of the revival disc is determined by the tidal torque, keeping t vis t GW . The mass accretion rate of the revival disc increases with time. In the late phase of the revival disc evolution, the mass accretion rate can exceed Eddington rate, and a relativistic jet is expected to be
Paper: 2020MNRAS.498.1319M (fulltext)
Wyatt ( 2014 ) showed that the disc should be truncated at a semimajor axis (at apocentre) where R H,Q is the Hill radius of an eccentric planet at apocentre. Under the influence of two planets, that inner edge could be further out if the strongest secular resonance is located beyond a in . Using equation (24) in Yelverton Kennedy ( 2018 ), we find the approximate location of the resonance to be (dashed blue line in Fig. 9 ) where a 1 and a 2 are the semimajor axes of the companions, while m 1 and m 2 are their masses, with 1 designating the innermost companion and 2 the BD. In order to effectively truncate the disc, the width of the secular resonance would need to be large enough to excite solids from a in to . This width is controlled primarily by the eccentricity of HD 206893 B. Given the uncertainties on its eccentricity () and its mass (12–50 M Jup ; Delorme et al. 2017 ), it is hard to assess whether secular resonances will deplete a region that otherwise would be stable. Assuming the best-fitting values of our fit and using equation (36) in Yelverton Kennedy ( 2018 ), we expect a secular resonance width of ∼20 au. Therefore, it is possible that the disc inner edge is not truncated by the BD, but by secular resonances. Note that HD 206893 B alone could truncate the disc out to ∼30 au if it has an eccentricity of 0.2 and a mass of 50 M Jup . A more precise characterization of HD 206893 B’s orbit and the disc inner edge could provide better estimates of the masses of the two inner companions. 5.2.3 Secular resonances between BD and planet at 30 au Secular resonances might still have created the gap if instead the resonance at 74 au was with the BD at 11 au and an outer planet sitting just interior to the disc inner edge at around 30 au. We find that this putative planet and the BD at 11 au would need to have a mass ratio close to unity (see solid orange line in Fig. 9 ). This is ruled out by direct imaging, which did not detect any companion more massive than 3 M Jup at 30 au (assuming an age of 50 Myr; Delorme et al. 2017 ). 5.2.4 Secular resonances between BD and planet at 74 au If the gap at 74 au was indeed carved by a planet at that distance, it is worth discussing what other observables could confirm this scenario. Given the range of possible masses of the BD and the putative planet c at 74 au, a secular resonance could be located in between the orbit of b and the disc outer edge. We find that the resonance would be located between 85 and 170 au for mass ratios of b and B of 10 −3 to 0.3 (solid green line in Fig. 9 ). This resonance would create a gap that is expected to be very wide (≳50 au if b has an eccentricity of ≳0.05 according to equation 36 in Yelverton Kennedy 2018 ), and thus noticeable by our observations. This means that we can already rule out that this gap is present in between 74 and 110 au, otherwise the gap would be seen to be much wider and the disc peak further out (or at 40 au instead). Therefore, imposing
Paper: 2018MNRAS.479.2415A (acknowledgments)
ACKNOWLEDGEMENTS ER, IFS and VL acknowledge the Programme National ‘Physique et Chimie du Milieu Interstellaire’ (PCMI) of CNRS/INSU with INC/INP co-funded by CEA and CNES. They also thank for generous financial support from La Région Haute-Normandie via the GRR Electronique, Energie et Matériaux, from the ‘Fédération de Recherche Energie, Propulsion, Environnement”, and from the LabEx EMC 3 and FEDER via the projects PicoLIBS (ANR-10-LABEX-09-01), EMoPlaF and CO 2 -VIRIDIS. IFS and VL thank PHC GALILEE 2018 PROJET (39379SF) and the GdR THEMS. IFS and JZM acknowledge support from the IAEA via the Coordinated Research Project ‘Light Element Atom, Molecule and Radical Behaviour in the Divertor and Edge Plasma Regions”. JZM acknowledges support from USPC via ENUMPP and Labex SEAM. This work is supported by BATTUTA Project (Building Academic Ties Towards Universities through Training Activities) in the frame of the Erasmus Mundus program, at LOMC UMR-CNRS-6294 of Le Havre University. YM thanks the SRI department, especially Mrs. Martine Currie, for outstanding hospitality.
Paper: 2016A&A...589A..53N (acknowledgments)
We are indebted to an anonymous referee, who provided very useful comments and suggestions from which we have indeed learned a lot, and which led to significant improvement of this paper. This work has been supported by Hungarian OTKA Grant NN 107637 (PI Vinko).
Paper: 2019ApJ...881...39H (acknowledgments)
We thank Astro Hack Week for providing the venue that started this project and the Lorentz Center workshop “The X-ray Spectral-Timing Revolution” (2016 February) that started the collaboration, as well as the Google Summer of Code Program for funding five coauthors (O.M., U.M.K., H.M., H.R., S.S.) who implemented a large fraction of the various library components over three summers. We thank Jim Davenport for providing help with the Kepler data. We are grateful to Matteo Guanazzi and Erik Kuulkers for their advice and suggestions for implementing DAVE, as well as the anonymous referee and the AAS statistics editor for their helpful comments on the manuscript. D.H. acknowledges support from the DIRAC Institute in the Department of Astronomy at the University of Washington. The DIRAC Institute is supported through generous gifts from the Charles and Lisa Simonyi Fund for Arts and Sciences and the Washington Research Foundation. M.B. is supported in part by the Italian Space Agency through agreement ASI-INAF n.2017-12-H.0 and ASI-INFN agreement n.2017-13-H.0. A.L.S. is supported by an NSF Astronomy and Astrophysics Postdoctoral Fellowship under award AST-1801792. E.M.R. acknowledges support from Conselho Nacional de Desenvolvimento Científico e Tecnológico (CNPq—Brazil).
Paper: 2017MNRAS.465.1889G (fulltext)
some insight into the spatial structure of the segregated groups. The upper right panel of Fig. 2 shows the map of the median positions of the segregated bins over the grey-scale density contours of the group A sample, using the same colour code as in the spectrum and histogram. As could also be seen in the contour plot from Fig. 1 , group A has a bilobed northsouth double peak structure. The two bins associated with segregated group (in yellow) are located towards the western medium density structure and the lower density peak. When we consider the RV, these bins are consistent with the unconnected components A2 and A3 previously described. The bins associated with group (red) are spread over the bilobed structure: the two that comprise 1 are associated with the southern and the one constituting 2 with the northern lobe. The bins associated with segregated groups (green) and (cyan) are both clearly associated with the upper density peak. In the lower right panel of Fig. 2 , the median RV of each bin is shown against its median (with segregated bins coloured with the same colour code as in previously described panels), to highlight the trend of the whole of group A: bins with higher RV tend to be located towards the north. The trend is noisy but clear, particularly when we consider a linear fit of the data, shown in Fig. 2 as a purple line. Fig. 3 shows in four panels the coloured contours of the density distribution of the segregated groups plotted over the grey-scale density of the complete sample, following the same colour code as in Fig. 2 : from left to right and top to bottom the density map of (yellow), (red), (green), and (cyan). Once again, the relationship of the RV of each group with the declination of its density peak is clear. Group is mostly concentrated around the lower peak and the east/central structure, with the highest density in the western area. Group , which is also the one with the most members, is distributed along the whole bilobed structure, reaching the highest density in the lower peak. Group is dense in both peaks, but mostly in the northern one, with a trend towards the east. Finally, group is very clearly associated with the northern peak of the bilobed structure. When we consider the unconnected components of group 1 and 2 the general RV declination trend still applies: group 1, with a lower RV, peaks in the lower lobe, while 2 peaks around the centre of the spatial domain, between both lobes. Figure 3. Spatial distribution of the four segregated groups from the analysis of group A. In each panel, the coloured contours show the density star distribution of each segregated group, with the same colour code as in Fig. 2 , over the grey-scale density of the group A. From left to right and top to bottom: group , in yellow, with kms 1 , group , in red, with kms 1 , group , in green, with kms 1 , and group , in cyan, with kms 1 . We will now compare these results with those obtained by Frsz etal. ( 2006 ) and Tobin
Paper: 2018MNRAS.474..760B (fulltext)
f NS , One can see from these equations that the luminosity of seed photons produced by the NS surface anticorrelates with the NS spin. For higher spin sources, supply of the seed photons will be reduced resulting in stronger Comptonization, in particular in large Compton amplification factor A and Comptonization parameter y , in the same way as the dichotomy between BH and NS sources had been explained in Paper I . Furthermore, assuming that the emission from the boundary layer is well approximated by a blackbody spectrum, then L bl ∝ T 4 . This means that if all other factors (especially the emitting surface area) are similar across the population of NSs, the temperature of the seed photons . The maximal Keplerian frequency in the accretion disc is of the order of f K ∼ 1.6–1.8 kHz (Sibgatullin Sunyaev 2000 ), while NS spins vary in the f NS ∼ 0.3–0.6 kHz range. At L X / L Edd ∼ 0.05, the seed photon temperature varies in the kT BB ∼ 0.8–1.0 keV range, i.e. by a factor of about ∼1.2–1.3, between low- and high-spin sources (Fig. 3 ). In order to explain such a difference between f NS ∼ 0.3 kHz and f NS ∼ 0.6 kHz sources, α ∼ 0.7 is required, which is broadly consistent with the conclusions of Paper I that the accreting material reaches the surface of the NS still possessing about ∼1/3–1/2 of its initial energy. 5.1 Can fast spinning NSs mimic BHs? From equation ( 2 ), at sufficiently high NS spin frequencies the luminosity of the boundary layer can become small. Theoretically, this can make accreting NSs similar to BHs from the point of view of the energy balance in the Comptonization region. Indeed, such a behaviour is evident in Fig. 4 where the high-spin sources at high Eddington ratios progressively overlap with the region of Comptonization parameters, identified with BHs in Paper I . This dilutes the boundary between BHs and NSs. From Fig. 2 , one can see that three sources expand into y ≳1.0 regions at higher L / L Edd : 4U 1636–536, KS 1731–260 and, most significantly, GS 1826–238. This behaviour is strongest at the highest Eddington ratios (but still compatible with the source being in the hard spectral state). All these sources have high spin (>500 Hz), and their behaviour markedly contrasts with that of the slower sources at similar L / L Edd , such as 4U 1728–34, the Comptonizing strength of which remains low. For GS 1826–238, which has one of the highest spins and mass accretion rates in our sample, the Comptonization parameter can reach values y ∼ 1.25, which is inside the region occupied by BH LMXBs. 2 However, such high values of y are achieved only at high Eddington ratios, L X / L Edd ∼ 0.1, at which NSs are characterized by extremely low electron temperature, kT e ∼ 10–20 keV (Fig. 2 ), well below values typical for BHs (fig. 5 in Paper I ). This behaviour in kT e − y distinguishes NSs from BHs and is illustrated in Fig. 5 . A detailed comparison of high-spin NSs with BHs will be a subject of a separate study. As a final
Paper: 2020MNRAS.492.1914N (fulltext)
is shown in Fig. 4(b) . Figure 4. (a) Reconstruction of B on a 1000-yr time-scale. Grey circles – 400-yr reconstruction. (b) Reconstruction of B on a 3200-yr time-scale. The grey areas are error intervals. 6 RESULTS AND DISCUSSION Here, using the MSR method, implemented by using equations ( 1 )–( 4 ), we have built reconstructions of the variations of the IMF strength B on different time-scales: 170, 400, 1000, and 3200 yr, based on different proxies ( IDV , Wolf numbers 2.0, and radiocarbon reconstruction). Typical correlation coefficients that we obtained are within the range 0.92–0.95, which on the one hand indicates the success of B reconstructions, and on the other hand, shows that we do not take into account all the factors of B variations in proxies, but only ∼85–90 per cent of the dispersion. Changes in IMF strength B vary significantly: for annual values , for a decade DB ∼ 2.5. Note that the cyclicity that makes the largest contribution to the dispersion in Fig. 4(b) is made by variations with the period slightly longer than 200 yr (Suess cycle). In order to see the most long-term trends in IMF variations, we compose the average values of B over 210 yr (see Fig. 5 ). Figure 5. The variations in the last 3200 yr of the B mean values for 210 yr (black circles) and their smoothed values over three points (white circles). Dashed line – linear trend, dotted line – mean, and dash–dotted line – forecast. Noteworthy is the long-term decreasing trend, which is −(0.39 ± 0.17) · 10 −3 nT· yr −1 over the past 3200 yr. Although its significance is ensured only by 2.3σ, the values smoothed over three points and the behaviour in time relative to the average (Fig. 5 ) speak of its reality. We also note that geomagnetic changes are unlikely to have contributed to the nature of this trend since we believe that the reconstruction used by Usoskin et al. ( 2014 ) is based on modern geomagnetic dipole moment data from the 1000 archeomagnetic field models (Licht et al. 2013 ). We also believe that the relationship between IDV and sunspot numbers depends only on the time-scale τ, but not on time t : any reconstruction is generally based on this assumption of time invariance. At the same time, an episode of increasing B has been observed over the past ∼400 yr. This means that the cosmic ray flux near the Earth’s orbit during these years decreases with time. As speculation, according to the Tinsley–Heelis mechanism (Tinsley Heelis 1993 ), this should lead to a decrease in the reflective properties (albedo) of the Earth’s atmosphere and an increase in the solar energy flux that reaches the Earth’s surface. All of this can make a partial contribution to the global warming effect. On the other hand, judging by the general nature of the variations in Fig. 5 , in the next 200–250 yr we will face an episode of a decrease in the IMF strength so that the increase in the Earth’s global temperature will decrease or decrease the increase rate if it is really related
Paper: 2015ApJ...809..152C (acknowledgments)
M.C. is supported by the Scientific and Technological Research Council of Turkey (TÜBİTAK) through the Research Fellowship Programme for International Researchers (2216). We thank Amy Y. Lien for valuable suggestions regarding the analysis of Swift BAT survey data.
Paper: 2018ApJ...853..143W (fulltext)
of horizontal field changes in shear region P is longer and followed by a flatter gradual increase. 4. Summary and Discussions We have presented a detailed study of the structural evolution of photospheric flow and magnetic fields associated with the 2015 June 22 M6.5 flare, concentrating on the central penumbrae of the δ region around the flaring PIL. We tracked the penumbral flow field with the DAVE method using high-resolution TiO images from BBSO/GST, and also analyzed the associated magnetic field using SDO /HMI vector magnetograms. The main results are summarized as follows. 1. Beginning from the first HXR peak, one segment of penumbra lying in the negative field experiences a distinct grow, extending northward at 1.5(±0.05) km s −1 for a distance of 26(±3); meanwhile, the flow velocity within the entire penumbral region becomes more pronounced, gaining a 115(±10)% increase reaching up to 2.2(±0.11) km s −1 at ∼90 minutes after first flare peak. These structure and flow field evolutions are accompanied by a step-like increase of horizontal magnetic field by 150(±15) G and decrease of inclination angle by 5(±0.5)° in 18 minutes from 17:50 UT to 18:08 UT. 2. A region of shear flow at the flaring PIL is found next to the location of a flare precursor brightening, and expands significantly after the flare. From the first HXR peak, the shear flow velocity increases gradually by 0.4(±0.1) km s −1 from 0.5(±0.05) km s −1 at 17:30 UT to 0.9(±0.05) km s −1 at 19:30 UT. As for magnetic field properties in the shear flow region, the horizontal field strength increases 220(±20) G while the inclination angle decreases 5(±0.5)° from 17:50 UT to 18:20 UT, both following a stepwise fashion. In comparison with the extending penumbral segment, the transition time of step-like magnetic field changes in the shear flow region is ∼10 minutes longer. Our results of the strengthening of penumbral structure at the center of the δ region, the associated increase (decrease) of horizontal field (inclination angle), and their close timing relationship to the flare energy release strongly favor the back-reaction of coronal fields as the underlying cause, which results in a more horizontal configuration of the photospheric field. While most other work only deals with the flare-related overall intensity change of penumbral structure, we are able to study the evolution of penumbral flow field owing to the high-resolution BBSO/GST data. We show that, although noticeable intensity darkening is only observed in the newly formed penumbral portion, the strengthening of penumbra is predominantly manifested as the enhancement of penumbral flow velocity and enlargement of penumbral area. The inconsistency of our results with the claim of darkening of the central penumbral feature in some previous studies may be due to the limitation of the low-resolution images used (e.g., Liu et al. 2005 ) and also the complexity of the central penumbral area near the PIL (Wang et al. 2017
Paper: 2018MNRAS.476.1140A (fulltext)
part of the H α line is not due to chromospheric activity. Figure 11. Variations in the Mg ii λ2800 emission line in several IUE spectra. [A colour version of this figure is available in the online version.] Following the procedure explained in Acker Jasniewicz ( 1990 ) for Abell 35, we have tried to find a modulation of the H α profile with the rotation phase, by measuring the flux variation between the nearby continuum and the deepest point of the H α absorption core. In contrast to the results found by Acker Jasniewicz ( 1990 ) in the case of Abell 35, we do not find any correlation with the rotational period during the long time span of the observations in LoTr 5. We have inspected the H α profile in all the Hermes spectra and we found no correlation with the rotation period, see Fig. A1 , where we show the phased-folded trailed spectra of the H α line with the rotation period. The Ca ii HK λλ3933, 3967 lines appear in emission in the HERMES data set (as already reported by Strassmeier et al. 1997 ). We do not find variations in these lines at the level of the H α variations that we find in HERMES data. Hence, Ca ii HK lines are not useful for studying the activity of the star. 3.5.2 Accretion disc around the possible close companion The presence of an accretion disc in LoTr 5 has been also discussed in the literature. If indeed the H α emission was due to an accretion disc, then the double-peaked profile would come from two different emission regions instead of a broad emission feature. Montez et al. ( 2010 ) already discussed the system within this scenario: since the presence of an accretion disc would imply the RLOF, they calculated the Roche lobe radius ( R L ) assuming the synchronization of the rotation and orbital period (i.e. they assumed a close 0.6 M ⊙ companion at ∼5.9 d). Under this configuration, the R L was below the radius of the star, filling its Roche Lobe and allowing the existence of an accretion disc. However, from our radial velocity analysis, we know that it is not possible to have such a massive companion so close to the G-type star. On the contrary, in the case we would have a low-mass object at ∼5.9 d, the R L would be ∼16 R⊙, and the star could not have filled its Roche lobe. Therefore, with the current data, we conclude that it is not possible to have an accretion disc via RLOF in LoTr 5. Finally, and as mentioned in Section 3.2 , no significant infrared excess is identified in the SED, which would point out that an accretion disc is highly unlikely in the system. 3.5.3 The possible resemblance with symbiotic stars Finally, it is worth mentioning that this double-peaked H α emission resembles the observed in some symbiotic stars (SS), the widest interacting binary systems known with a white dwarf accreting material from a giant star via stellar winds (see e.g. Van Winckel, Duerbeck Schwarz 1993 ; Burmeister Leedjärv 2009 ). These systems are sometimes embedded into a bipolar nebula similar to some PNe (like