adsabs/WIESP2022-NER huggingface dataset:
Paper: 2017MNRAS.468.1236B (fulltext)
includes several vestoids and inner belt non-vestoids, while the second group includes some inner belt vestoids and non-vestoids, as well as middle/outer belt non-vestoids. It is worth noting that the errors involved in the determination of spectroscopic and polarimetric parameters are usually large. In many cases, the differences observed amongst the samples fall within their 1 uncertainties, making them indistinguishable in practice. In principle, there is no clear relation between these spectroscopic and polarimetric differences and the different dynamical populations. 2.3 Size distribution So far, the clearest evidence that prevents a link between (4)Vesta and (1459)Magnya is the size of the latter, which does not fit within the expected size frequency distribution (SFD) of the ejecta from craterization events (Durda etal. 2007 ; Michel etal. 2015 ) like those that excavated the surface of (4)Vesta. On the other hand, the remaining non-vestoids (from both the spectroscopic and SDSS samples) are all as small as the vestoids and may also have originated from craterization events. Images of Vestas surface taken by the Dawn probe revealed two large basins in the south pole of the asteroid. One of them, named Rheasilvia with 500km of diameter, has been dated 1Gyr ago. The other one, named Veneneia with 400km of diameter, has been dated 23Gyr ago (Marchi etal. 2012 ). While Rheasilvia is the likely source of the current vestoids (Buratti etal. 2013 ; McSween etal. 2013 ), Veneneia could be the source of several non-vestoids (Schenk etal. 2012 ). Pirani Turrini ( 2016 ) studied the formation of large basins on the surface of (4)Vesta during the epoch of the giant planet instability and concluded that any record of such basins would have been erased by the subsequent crater saturation over the age of the Solar system. This is in line with the fact that observed craters on Vestas surface dated less than 4.1Gyr ago (OBrien etal. 2014 ). Therefore, the possibility of a very early craterization event on the surface of (4)Vesta that led to the origin of a Vesta-like family cannot be ruled out. In this context, it is interesting to discuss the SFD of the different dynamical populations of V-type asteroids. This is shown in Fig. 2 for the sample of SDSS V-type candidates, but a similar result was obtained by Nesvorn etal. ( 2008 ) considering the spectroscopic sample. The fugitives population shows a similar SFD to the vestoids, indicating that these two populations may be related and have the same collisional age. On the other hand, the low-inclination population and the middle/outer population show SFDs that differ from the Vesta family, suggesting that these populations have a different origin/age. These populations (excluding (1459)Magnya due to its size) could have been originated by early craterization events on the surface of (4)Vesta, like the Veneneia basin or even older events (Schenk etal. 2012 ; McSween etal. 2013 ). Figure 2. The cumulative
Paper: 2016MNRAS.458.3904R (acknowledgments)
We thank an anonymous referee for some very helpful and insightful comments. AV is supported by the National Science Foundation Graduate Research Fellowship, Grant No. DGE 1144152. TK thanks Cheryl Healy and Beth Christie for material support in telescope operations and equipment.
Paper: 2017ApJ...841...86B (acknowledgments)
E.V.Q.’s research was partially supported by a NASA Postdoctoral Program Senior Fellowship at the NASA Ames Research Center, administered by Universities Space Research Association under contract with NASA. The simulations presented here were performed using the Pleiades Supercomputer provided by the NASA High-End Computing (HEC) Program through the NASA Advanced Supercomputing (NAS) Division at Ames Research Center. The authors would like to thank Chris Henze for his expertise and assistance in performing simulations on the Pleiades Supercomputer.
Paper: 2016ApJ...824...74G (fulltext)
I -band model at 1% significance. This could be due to blending problems and also to the lack of extinction correction in OGLE-IV data. To perform stability tests, the data are split by dummy variables into two data sets around P = 3.85 days (OGLE-II) and P = 3.5877 days (OGLE-IV), so as to have about half of observations in each set. No evidence of structural breaks is found, because the P -values obtained, by testing versus the hypothesis that at least one of them is not 0 in Equation ( 32 ), are 0.601 (OGLE-II) and 0.453 (OGLE-IV) in the V -band, and 0.810 (OGLE-II) and 0.267 (OGLE-IV) in the I -band. Models explain about 61% and 68% of variability in the V - and I -bands respectively (OGLE-II), and 54% and 65% of variability in the V - and I -bands respectively (OGLE-IV). The estimated parameters of the models and their confidence intervals are reported in Table 6 . Results for the I -band are shown only for information purposes because of their specification problems. Table 6 Optical Leavitt Law for the LMC Note. The reported β estimators are those given by Equations ( 26 ) and ( 27 ). They were obtained by using the MM -regression. Confidence intervals are reported at 95%. The last columns report the values of zero point and the slope obtained by Soszyński et al. ( 2015 ) using the sigma-clipping algorithm and OLS over OGLE-IV data. The OGLE-II and OGLE-IV MM -estimators obtained in the V -band and their respective confidence intervals include at the OLS estimators reported by Udalski ( 2000 ) and Soszyński et al. ( 2015 ). The OGLE-IV LMC data were also split into two around P = 10 days by Equation ( 32 ). No evidence of a structural break is found in the slope, because the P -value obtained is 0.705. However, the zero-point adjustment is significant. As a result, there is evidence of two parallel regression lines. The SMC slope tests for models ( 26 ) and ( 27 ) are significant because the P -values obtained are less than 0.0001 for the V - and I -bands (OGLE-II, OGLE-IV). However, these models show problems. The test of Godfrey and Orme finds enough evidence of specification problems because the P -values obtained are less than 0.0001 for the V - and I -bands (OGLE-II, OGLE-IV). To perform stability tests, data are split around 2.0039 days (OGLE-II) and 1.9429 days (OGLE-IV) so as to have about half of observations in each data set. Enough evidence of structural breaks is found: P -values are less than 0.0001 in both bands (OGLE-II) when testing versus the hypothesis that at least one of them is not 0 in Equation ( 32 ). Besides, when individual significance tests are performed for and , the P -values obtained are less than 0.0001 for in both bands, and 0.0322 and 0.0082 for in the V -band and I -band respectively. As a result, in both bands there are two regression lines at 5% significance (OGLE-II). Unlike OGLE-II models for the SMC, there is not enough evidence to support the idea that OGLE-IV model parameters for the slope are
Paper: 2019ApJ...887...91M (fulltext)
with black open circles. We warn the reader that we cannot rule out the presence of other binaries with no obvious s -elements and/or RV differences among our remaining 15 giants. Another stellar population that may affect the appearance of the ChMs and likely has a connection to binary stars is the blue stragglers (BSs). In particular, we consider their evolved counterparts that lie close to the RGB. In Figure 7 , we use blue symbols to represent the candidate BS stars of NGC 3201 in the CMD. The objects marked with blue asterisks are possible BSs selected on the CMD that are evolving toward the RGB phase, which clearly displays lower values in the diagram (top right panel). The BSs in GCs are tightly linked to the cluster binary populations (Knigge et al. 2009 ). The BSs in open clusters themselves have a very high binary fraction (Mathieu Geller 2009 ), and preliminary evidence suggests the same is true in NGC 3201 (Giesers et al. 2019 ). When we include the evolved BSs in the ChM (bottom panel of Figure 7 ), they populate a well-defined sequence on the blue extension of the 1G sequence. We conclude that, based on RVs and s -process element abundances, the three bluest 1G stars that we analyzed spectroscopically are binary systems. Some of them are possibly associated with the BS population of NGC 3201, as evolved BSs would contribute to the bluest extension of the 1G sequence in the ChM. Figure 7. Top panels: CMD (left) and diagram (right) of NGC 3201. Bottom panel: ChM of NGC 3201, similar to Figure 1 but extended to show the evolved BSs. In all panels, the green dots represent our analyzed spectroscopic targets, the open blue circles are the BSs, and the blue asterisks are the evolved BSs. 6.1. Noninteracting Binary Simulations To investigate the effect of noninteracting binaries formed by pairs of 1G stars, we reproduce in Figure 8 five isochrones, I1—I5, with [Fe/H] = −1.50, [ α /Fe] = 0.4, and an age of 13 Gyr from Milone et al. ( 2018 ). These isochrones are constructed with a different combination of He, C, N, and O to represent the typical chemical pattern of different stellar populations in GCs. We used the isochrone I1, which correponds to the 1G, to generate a stellar population of 1G–1G binaries that we represent with black points in Figure 8 . Figure 8. Left panels: simulated CMD, with five isochrones corresponding to different stellar populations with different He, C, N, and O. Here I1 is the isochrone corresponding to the first population; I2 has the same abundance of C, N, and O as I1 but different He; and in I3, I4, and I5, the abundances of He, C, N, and O have been varied to represent the second population chemical pattern. Right panel: location of the five different stellar simulated populations at a reference magnitude on the ChM plane. The black points are noninteracting binaries formed by pairs of I1–I1 stars. Binaries can contribute to the width of the 1G in the ChM, but only binaries with a large mass ratio ( q ≳
Paper: 2016MNRAS.457.4285S (acknowledgments)
We thank the anonymous referee for useful discussion.
Paper: 2018ApJ...865..124M (fulltext)
cool via synchrotron and IC losses. Following Murase ( 2012 ), the deflection angle during the radiation cooling time is given bywhere t syn is the synchrotron cooling time, r L is the Larmor radius, and the Lorentz factor of pairs in the black hole rest frame is . Only synchrotron cooling is considered, because the IC cooling is suppressed for such high-energy electrons and positrons. The above equation implies that the pairs lose their energy after they become isotropized. Note that the X-ray emission is doubly suppressed, because the deflection is larger than the jet opening angle, i.e., , and the time spread in the cascade emission is longer than the intrinsic flare duration, i.e., . Also, the CR neutrons do not produce extra pairs via the Bethe–Heitler process. In addition, in the np scenario, X-rays could also be isotropized via the Compton scattering. Thus, X-rays from the neutral beam can be highly suppressed if the CR energy is lower than UHE energies, and the associated hadronic cascade signatures can be masked by the GeV γ -ray emission from the blazar zone. Note that the UHE neutron beam, if it exists, can also produce beamed cascade emission at TeV energies. The resulting synchrotron pair-echo emission (Dermer et al. 2012 ; Murase 2012 ) is expected to have some time delay. We also remark that the very high energy flare seen by MAGIC was detected about 10 days later after IceCube-170922A. 4. Summary We considered implications of the high-energy neutrino flares from TXS 0506+056 by examining various constraints. Flaring blazars could be the brightest sources in the neutrino sky, while the observations can most naturally be reconciled with existing theoretical models, if the blazars are subdominant (∼1%–10%) in the diffuse neutrino intensity. Interestingly, within the standard leptonic scenario of γ -ray emission, we found that the blazar neutrino emission itself can readily be dominated by flaring episodes. This could explain why the significant neutrino signals from TXS 0506+056 were found only in the time-dependent search and that the steady component that typically dominates γ -ray emission has not been seen yet in neutrinos. Bright neutrino flares like the ones observed for TXS 0506+056 could contribute only up to a few percent of the INB. If the association with this blazar is physical, such flares can be detected with a rate of by dedicated time-dependent searches in the near future. Based on analytical considerations, we also showed the importance of X-ray constraints to test the physical models of TXS 0506+056. An efficient electromagnetic cascade is unavoidable in the canonical blazar models based on the single radiation zone, which also predicts that the 2014–2015 neutrino flare found in the archival data should be accompanied by X-ray emission with . This can readily be tested by X-ray sky monitors such as Swift and MAXI . Nondetection of the X-ray flares will necessitate more complicated models involving multizone
Paper: 2017ApJ...850..142C (acknowledgments)
It is a pleasure to acknowledge fruitful discussions with Fatima Ebrahimi, Eero Hirvijoki, Hantao Ji, Russell Kulsrud, Roscoe White, and Yao Zhou. This research was supported by NSF Grant Nos. AGS-1338944 and AGS-1460169 and by DOE Grant No. DE-AC02-09CH-11466.
Paper: 2016A&A...587A..70S (fulltext)
drop in the abundances, which is and . A KS test suggested that the two distributions are different (-value of 3%), consistent with a radial movement of the gas as the cause of this feature. Finally, of the galaxies displaying this inner abundance drop neither showed evidence of a star-forming ring nor of a bar. This suggests that another mechanism related to radial movements is responsible for causing this feature. To shed some further light on this question, we investigated the presence of this feature depending on the integrated stellar mass of the galaxies. We conclude that the galaxies displaying the strongest oxygen abundance inner drop are the most massive ones (Fig. 9 ), suggesting that stellar mass plays a key role in shaping the inner abundance profiles. In summary, this is the first study that analysed the oxygen abundance distribution for a sample of face-on spiral galaxies spaxel by spaxel. Our final sample of 122 galaxies provided more than 185 000 oxygen abundance values (8200 of them beyond two disc effective radii) to carry out the analysis. The results were compared to those obtained following the classical procedure of detecting H ii regions, leading to equivalent results that point to the same conclusions: (i) the existence of a common abundance gradient, independent of other properties of galaxies, in particular the presence of bars, which seems not to have the flattening effect predicted by numerical simulations; (ii) the existence of a flattening of the abundance gradient in the outer regions of discs, which seems to be a common property of disc galaxies; and (iii) the existence of a drop of the abundance in the inner regions of disc galaxies, only visible in the most massive ones and most probably associated to radial movements of the gas (sometimes a bar or a circumnuclear star-forming ring). These results support the scenario in which disc galaxies present an overall inside-out growth. However, clear deviations were shown with respect to this simple scenario that affect the abundance profiles in both the innermost and outermost regions of galaxies. The agreement between the two methods arises because the spaxel size of the CALIFA datacubes is of the order of the size of a typical H ii region. However, this is not expected for IFS data with better spatial resolution and a smaller spaxel size, where different areas of the H ii regions can be resolved and even abundance gradients can be found. Different procedures that can accomodate data that are capable of resolving H ii regions may be needed. Appendix A: Fundamental parameters and oxygen abundance information In this section we present a table with general information and derived oxygen abundance information for all the galaxies in the sample. The meaning of each column is indicated in the notes below. Table A.1 Fundamental properties and oxygen abundance information. Notes. (a) The galaxy name; (b) the morphological type; (c) -band absolute magnitude; (d) the
Paper: 2020MNRAS.494..378D (fulltext)
used the local Maxwellian approximation to generate their ICs, which can lead to artificial relaxation in the centre of the subhalo (Kazantzidis et al. 2004 ). Overall, it appears that the fits from Green van den Bosch ( 2019 ) perform slightly better than all other models, but our model is at least comparable to, if not more accurate than the models presented in Hayashi et al. ( 2003 ) and Peñarrubia et al. ( 2010 ). Figure 1. Comparison of different model descriptions of the density, enclosed mass, and circular velocity profiles of the bound satellite remnant after one orbit in the Slow (left-hand panel) and Fast Simulation (right-hand panel). Simulation results are shown with points, and the models are shown with lines. The thick dashed curve shows the original, untruncated NFW profile. The residuals in the fits are calculated as ( y model / y simulation − 1), where y is the density, mass, or circular velocity profile. Our model from Paper I is described in Section 2 , and the other models in Appendix B . In all cases, the models are dependent on one parameter, the bound mass fraction, which we calculate as the self-bound mass in the simulation, as described in the text. We note that since all of these methods are dependent on the bound mass, they are sensitive to the method used for defining bound particles. Also, bound mass calculations include mass that is only temporarily bound, but is leaving the system (e.g. Peñarrubia et al. 2009 ). Most of this temporarily bound mass is at larger radii, which may explain why the mass-loss models all fail to match the bound mass profiles from the simulations beyond 4–5 r unit in Fig. 1 . A more accurate normalization procedure for our model might be to fit the radius enclosing some fraction of the mass, or to fit the peak circular velocity; rather than take either of these approaches, we will later use the mass-loss model in Section 6 to predict theoretically. 5 OUR MODEL ASSUMPTIONS The truncation method described in Section 2 determines which particles will remain bound, and which will be ejected from the system, given a tidal truncation energy . As shown in the previous section, this approach predicts the density (or equivalently the mass or the circular velocity) profiles of the remnants quite well, but appears to be more accurate when mass-loss is slower. Our goal in this section is to understand why this is the case, focusing on the Fast and Slow Simulations described in Section 3 . Suppose we consider a simple model where mass-loss is primarily an energy-based process, as proposed by Choi et al. ( 2009 ). Three main simplifying assumptions are: Mass-loss occurs episodically around the pericentre of the orbit; after each episode the system returns to equilibrium, with the modified spherically symmetric, isotropic distribution function predicted by the energy truncation method. Individual particles are lost or retained based solely on their energy with respect to the rest of the
Paper: 2018ApJ...863..177F (acknowledgments)
We are grateful to Eugene Churazov, Sergey Komarov, and an anonymous referee for insightful comments and suggestions. This work was supported partially by Natural Science Foundation of China (grant Nos. 11633006, 11725312, 11643001, U1431228, 11233003, 11421303, and 11703022), Natural Science Foundation of Shanghai (No. 18ZR1447100), and Chinese Academy of Sciences through the Hundred Talents Program and the Key Research Program of Frontier Sciences (Nos. QYZDB-SSW-SYS033 and QYZDJ-SSW-SYS008). G.M. acknowledges the support by the Fundamental Research Funds for the Central Universities (WK2030220017). X.E.F. thanks Zhen-Yi Cai for the help on scientific plotting. Some simulations presented in this work were performed using the high-performance computing resources in the Core Facility for Advanced Research Computing at Shanghai Astronomical Observatory.
Paper: 2018ApJ...861..154S (fulltext)
there are ≈100 M ⊙ of hot gas in the X-ray-emitting bright regions (see Table 3 ). There have been a variety of simulations that predict the amount and isotopic composition of debris from SN explosions. For example, Seitenzahl et al. ( 2013 , hereafter Sei ) list 17 models for a Type Ia explosion of a 1.4 M ⊙ white dwarf, and Woosley Weaver ( 1995 , hereafter WW ) simulate the gravitational collapse of 78 stars with zero-age main-sequence mass ranging from 11 to 40 M ⊙ s and metallicity from zero to solar values. Table 6 lists a few selections of predictions from these references. The WW examples have been chosen because of higher values for the amount of Si ejecta and to show a general trend. Table 6 LMC Abundances and Ejecta Mass The abundances of elements in the LMC ISM have been taken from Schenck et al. ( 2016 ), who give values for material swept up by the brighter LMC SNRs. Table 6 compares our abundance values with these. Although our uncertainties are not small, the elements O, Ne, Mg, and Fe appear to be present in amounts expected from the ISM. The only enhancement measured by the Chandra ACIS is the abundance of Si, which is 1–2× solar. Table 6 lists the Schenck abundances, our ACIS abundances, the mass of each element in 100 solar masses of material with solar abundances, the masses in 100 solar masses of material with LMC ISM abundances, the masses in 100 solar masses of X-ray-emitting material in J0453-6655 as determined by ACIS, and, in the last four rows, masses predicted by four models of SN explosions. In the 100 solar masses of diffuse gas we find ≈0.13 ± 0.05 M ⊙ of Si. Values for other elements are close to those expected for the ISM. Uncertainties are listed in Table 6 only for the element Si, but Table 2 gives all uncertainties in the spectral fits. There is good agreement between ACIS and EPIC results, but there is a dependence on model used, which could be a factor of 2. This is not important for the conclusions to be drawn here. The elements produced in an SN Ia are predominantly Fe (the 56 Ni produced quickly decays to 56 Fe), O, Si, and S. Almost all models yield several tenths of a solar mass of Fe. Our ACIS analysis finds 10× less, including the ISM contribution. Based solely on the lack of Fe, an SN Ia is an unlikely origin. An SN II, from gravitational collapse of a massive star, will have dispersed several solar masses of ejecta, and the amount of material ejected increases with stellar mass. The Si/O ratio generally decreases with increasing pre-SN star mass, but there is considerable variation from model to model. The predominant element ejected in the WW core-collapse simulations is O. All stars with mass above 20 M ⊙ and solar metallicity eject more than 1 M ⊙ of O, well above the ACIS measurement of 0.11 M ⊙ , indicating that the pre-SN star almost certainly is 20 M ⊙ in size. The data also favor an explosion that ejects at least 0.07 M ⊙ of Si and less than 0.1 M ⊙ of O. None of the WW simulations fit