adsabs/WIESP2022-NER huggingface dataset:
Paper: 2020ApJ...896..160L (fulltext)
degrees of freedom. The nonrelativistic Model 1 already gives a good fit to the X-ray spectra of Fairall 9. In addition, we do not clearly detect the signature of a broad iron line in Figures 1 and 2 . All these considerations lead us to conclude that the relativistic reflection component in the spectra of Fairall 9, if present, is weak. To constrain possible deviation from Kerr spacetime around the black hole in Fairall 9, we replace the component relxill in Model 3 with its non-Kerr extension relxill _ nk v1.3.2 (Bambi et al. 2017 ; Abdikamalov et al. 2019a ). Similar to the results with relxill , we also find three solutions. The best-fit parameters are shown in Table 5 and the constraints of the deformation parameters in Figure 6 . We only show the results corresponding to solution 1 (Model 3.1) because the deformation parameters are totally unconstrained in the other two solutions. This is easy to understand since both negative spin and truncated disk result in a large distance between the black hole and the inner edge of the accretion disk. In this case, the effect on the reflected spectrum due to deviations from the Kerr background are not strong enough to be detected with current data. When assuming solution 1, we find from Figure 6 that the deformation parameters can recover 0, which means the Kerr solution is recovered. Figure 6. Confidence contour plot for spin parameter and deformation parameters. The red, green, and blue lines correspond to 68%, 90%, and 99% confidence levels respectively. The gray region is not included in our analysis because the spacetime there has pathological properties. Table 5 Best-fit Values with Model relxill_nk Note. Model: Tbabs ∗ (nthComp + relxill _ nk + xillver + zCutoffpl) . The parameters with * are fixed at given values in the fit. Uncertainties are given at the 90% confidence level. 4. Discussion and Conclusions Fairall 9 has a number of interesting properties (listed at the beginning of this paper) that make it a promising source for testing the Kerr black hole hypothesis using X-ray reflection spectroscopy. Unfortunately, the relativistic reflection component, if present, is weak, which makes the selection of the correct astrophysical model with the available data challenging. However, we should stress that the ideal source with all the nice properties for testing general relativity and no bad properties probably does not exist, so we have to try to do our best with the available sources and data. Moreover, in the case of Fairall 9 we can expect (see Section 4.3 below) that future observations with a higher energy resolution near the iron line can solve the degeneracy and be capable of selecting the astrophysical model and the physical best-fit solution. 4.1. General Considerations We studied a simultaneous XMM+NuSTAR observation of Fairall 9. The soft excess in the spectra can be described by a Comptonization component (Model 1) and the fit is already good. When we add a blurred reflection
Paper: 2018MNRAS.477.3910B (fulltext)
taking its place. It is only when m 3 ∼ m 2 that genuinely chaotic dynamics can take place, in some cases leading to coalescence. Finally, Fig. 5 shows the merger fraction as a function of the initial eccentricity of the inner and outer binary. We note that the merger fraction increases with increasing e in , while it decreases with increasing e out . The dependence upon e in is readily understood, since highly eccentric inner binaries are closer to the efficient GW-emission stage and can easily be driven to coalescence by a relatively mild perturbation from a third body. The dependence upon e out is likely due to the fact that quasi-circular outer binaries form a stable hierarchical triplet for a comparatively longer time during the inspiral of m 3 , hence leaving more room to the development of K–L resonances, which are efficient at driving the inner binary to coalescence. Conversely, in very eccentric outer binaries, m 3 soon interacts with the inner binary at pericentre, entering the chaotic phase. Chaotic interactions are more likely to result in ejections rather than mergers, hence suppressing the overall merger fraction. Figure 5. Merger fraction (colour coded) as a function of the initial inner ( e in ) and outer ( e out ) binary eccentricities. 3.1.1 Importance of PN corrections As pointed out in the previous section, the K–L mechanism can be suppressed by general relativistic effects. In particular, relativistic precession tends to destroy the coherent pile-up of the perturbation that the third body induces on the inner binary, hence effectively damping the K–L resonances. 5 In order to quantify the impact of the relativistic precession on the merger fraction and to compare our results with previous work that neglected this effect (e.g. Iwasawa, Funato Makino 2006 ; Hoffman Loeb 2007 ), in Table 3 and in Fig. 6 we compare, only for the case m 1 = 10 9 M ⊙ , the merger fraction obtained with and without 1PN corrections. Overall, the merger fraction is substantially higher in the case without 1PN terms (right-hand panel of Fig. 6 ). As can be seen, the largest differences compared to the full case occur for q in ≪ 1, where, because of octupole-order terms, the K–L mechanism is maximally effective. On the contrary, for large q in the merger fractions with or without 1PN corrections are similar, because, as previously discussed, coalescences are mainly due to chaotic strong encounters, rather than to K–L oscillations. Our results highlight the importance of K–L resonances in inducing MBH mergers in triple systems, and the need to account at least for 1PN corrections. Figure 6. Comparison between the merger fraction in the case m 1 = 10 9 M ⊙ , when the dynamics is evolved with (left-hand panel) and without (right-hand panel) 1PN term. The suppression of the merger fraction for low q in , due to 1PN precession, is clearly visible (see discussion in the main text). Table 3. Comparison of the merger fraction from the simulations with m 1 =
Paper: 2018ApJ...863...74K (fulltext)
tips of the H 2 bow shocks, where the shock velocities are high and the gas will be dissociated, rather than behind the H 2 bow shocks (e.g., Davis et al. 1999 , 2000 ; López et al. 2010 ; Lee et al. 2014 ; Bally et al. 2015 ). Detection of further [Fe ii ] emission in G53.11_MM1 with high signal-to-noise ratio requires deeper imaging observations, although the marginal detection of [Fe ii ] emission can be interpreted as intrinsically fainter or absent [Fe ii ] emission compared to H 2 emission. In the outflows from high-mass YSOs, the [Fe ii ] detection rate with respect to the H 2 detection rate tends to be low, ≲50% or much less (Cooper et al. 2013 ; Wolf-Chase et al. 2013 ; Caratti o Garatti et al. 2015 ). The brightness of the [Fe ii ] lines also tend to be weaker than the brightness of the H 2 lines (Caratti o Garatti et al. 2015 ). This can be attributed to the different extinction effect between the two lines but it is more likely to happen because H 2 and [Fe ii ] emission arise from different physical conditions, such as gas density, temperature, or shock velocity. In the G53.11_MM1 outflow, the strong, extended H 2 emission with very weak or negligible [Fe ii ] emission may imply that slow, C-type shocks are dominant in G53.11_MM1. Further spectroscopic observations will be necessary to derive the physical conditions of the environment and confirm the shock’s properties. 5. Central YSOs The core G53.11_MM1 is bright from IR to millimeter, with a large IR-excess emission. The central star (YSO1) was previously identified in the Red MSX Source survey with a bolometric luminosity of (3–4) × 10 3 L ⊙ (G053.1417+00.0705; Mottram et al. 2011 ; Lumsden et al. 2013 ); however, it is in fact composed of two sources, YSO1 and YSO2, which are separated by ∼8′′ in the Spitzer mid-IR images with higher spatial resolution. Both YSOs are saturated in the MIPS 24 μ m image but their SEDs with strong excess in mid-IR and spectral indices derived by using the available photometry from the GLIMPSE and MSX catalogs classify them as Class I YSOs that have a dusty envelope infalling onto a central protostar (Kim et al. 2015 ). The two YSOs are also observed in near-IR wavebands. Both are fairly bright in the K -band, marginally detected in the H -band and are not observed in the J -band, which indicates that they are deeply embedded. The evolutionary stages of the YSOs, with the proximity to the center of the outflow (see Figure 2 ), suggest that one or both can be the driving source of the G53.11_MM1 outflow. The coordinates of YSO1 and YSO2 are ( α 2000 , δ 2000 ) = (19 h 29 m 1760, +17°56′233) and ( α 2000 , δ 2000 ) = (19 h 29 m 1726, +17°56′173), respectively. In the following, we will discuss their photometric variability in near-IR and physical parameters constrained from SED analysis. 5.1. Near-IR Photometric Variability YSOs are known to commonly show variability (e.g., Carpenter et al. 2001 , 2002 ; Morales-Calderón et al. 2011 ; Johnstone et
Paper: 2020A&A...644A..49M (fulltext)
as well (see, e.g., Gromadzki et al. 2013 , for more details about light curves of S-type symbiotics). 3.4. Hot component and outburst activity The most prominent features in the light curves of Gaia18aen (Fig. 7 ) are the outbursts observed in 2018. In Fig. 8 , the part of the light curve showing the active stage is depicted. For clarity, the light curves in various filters were shifted to the same level to study the structure of the active stage. Individual brightenings are labeled by the numbers 0−5. Fig. 7. Light curves of Gaia18aen. Individual light curves in various filters were shifted to the level in the OGLE V filter for clarity; values of shifts are shown in the figure legend. Different colors denote different filters. Fig. 8. Outbursts of Gaia18aen. The colors are the same as in Fig. 7 . The vertical dotted lines indicate the individual brightenings (denoted by numbers 0−5). The arrows show the times when the two spectra were obtained. Fig. 9. Light curves in selected filters phased with the period of P = 486.9 d. Individual datasets were shifted to the level in the OGLE I filter. The first small pre-outburst of about 0.5 mag in V filter (denoted 0) was detected at JD 2 458 089. It was followed by more prominent outburst (1) after approximately 27 d (JD 2 458 116) of gradual decline. The maximal brightness during this outburst was ≈11.2 mag in V filter (JD 2 458 119), thus the amplitude was about 3.3 mag in comparison with the quiescent V magnitude (OGLE V ∼ 14.5). The analysis of the combined light curves constructed from the observations in various filters revealed the complex structure of this outburst with at least two another re-brightnenings that followed the first maximum. Approximately after 47 d (JD 2 458 163) after the first maximum Gaia18aen reached the magnitude similar to the brightness of the pre-outburst 0. The outburst denoted in this paper as 2 started after 57 d of a quasi-steady period (JD 2 458 220) and lasted for 35 d (until JD 2 458 255). The maximal brightness in V filter during this period was ≈12.5 mag. A third, small-scale outburst was detected at JD 2 458 357. The amplitude of the outburst was much smaller than in the previous cases, however, the approximate duration was similar to the outburst 2. Up-to-now, the last two increases in brightness (denoted 4 and 5) were detected at JD 2 458 470 and JD 2 458 503. It is worth noting that both brightenings were much shorter than the previous events, at 12 and 15 d, respectively. Moreover, the shape of at least the first event (4) looks as thought it was suddenly interrupted, which may indicate that both brightenings are part of a single outburst that is apparently interrupted by some obscuration process. Since JD 2 458 518, there were no other significant brightenings detected in the case of Gaia18aen. The amplitude of the outbursts and their duration resemble the behavior of typical classical symbiotic stars (e.g., AG Dra, Z And; see, e.g., Mikołajewska et al.
Paper: 2016ApJ...832..119H (fulltext)
galaxies. These objects were used to draw conclusions about the AGN population that would be missed in optical and X-ray surveys. Of the 30 AGN candidates using the Jarrett et al. ( 2011 ) selection box and presented in Satyapal et al. ( 2014 ), only one object is in the mass and redshift range probed by our analysis here (the others are either at higher masses or higher redshifts), and it is found in our sample of star-forming dwarfs in the Jarrett et al. ( 2011 ) box (NSA ID 93798). Secrest et al. ( 2015 ) measured an X-ray luminosity from XMM-Newton observations (extracted with a ∼15′′ aperture) of L 2–10 keV = 2. 4 × 10 40 erg s −1 for this object, and claim that the W1 − W2 color is strongly indicative of AGN activity. While this object is indeed in the Jarrett et al. ( 2011 ) selection box, given the median uncertainties on the WISE photometry, its infrared colors are consistent with this object potentially being scattered from the primary locus of star-forming galaxies and it does not have a strong power law to longer infrared wavelengths as expected for an AGN. Sartori et al. ( 2015 ) presented a sample of 83 objects with photometry placing them in the Jarrett et al. ( 2011 ) selection box, using data from the older WISE All-sky data release. These authors only required S/N > 2 for detections in the W1 , W2 , and W3 photometric bands. We cross-matched against the more up-to-date AllWISE selection, and using a S/N cut of 3.0, we find that only 34 objects (∼41%) continue to fall inside the Jarrett et al. selection box. 9 9 All 83 of the objects in the Sartori et al. ( 2015 ) sample have updated AllWISE W1 and W2 photometry with S/N > 3.0, but 26 of the objects have W3 photometry with S/N 3, only 10 of which have 2 S/N 3. Of these 34 objects, only 15 are in our parent sample of dwarf galaxies. The remaining 19 objects are either not in the NSA because their redshifts are greater than 0.055, or are in the NSA but not classified as dwarf galaxies by our mass criterion. Looking closer at these 15 dwarf galaxies, 10 are found in Table 1 (NSA IDs 4610, 6205, 10311, 53934, 57649, 93798, 109919, 117162, 118025, and 118523), 4 of the 15 are BPT-selected AGNs or composite galaxies (NSA IDs 10779, 79874, 124050, and 125318), and one is the previously discussed quasar contaminant (NSA ID 110100). One object that was listed as a Stern et al. ( 2012 ) candidate but not a Jarrett et al. ( 2011 ) candidate in Sartori et al. ( 2015 ) moved into the Jarrett et al. ( 2011 ) selection box with the updated AllWISE photometry (NSA ID 98135), and one object we have listed in Table 1 does not appear in Sartori et al. ( 2015 ) (NSA ID 151888). 3.5. The Feasibility of Detecting Supermassive Black Holes in Dwarf Galaxies with WISE Finding evidence for AGNs in dwarf galaxies is challenging as AGNs at a given accretion rate have lower bolometric luminosities due to the smaller masses of the central black holes, making them harder to observe above the emission due
Paper: 2018ApJ...865..142B (fulltext)
It is defined as the standard deviation of the differential light curve over each of the five rotations (Mathur et al. 2014 ). In the analyses below, all three of these measures were tried and found to be essentially equivalent for the purpose of correlating brightness with variability. The method actually employed in this paper is the re-sampled fit to the smoothed R loc ; we will call this R smo . This leads to a set of both absolute brightness and variability measurements at the same 45 time points in each light curve. A few illustrations of these sets are provided in the upper panel of Figure 11 . A linear fit is made to each set of points, and both the slope and Pearson-r correlation coefficient are computed. For the small and large red asterisks, the [slope, Pearson-r] values are [−0.14, −0.91] and [−0.92, −0.84]; for the small and large pluses, they are [0.16, 0.87] and [0.59, 0.53]; and for the small and large black diamonds, they are [0.00, 0.01] and [−0.20, −0.30]. The KIC numbers for the stars shown (in the same order) are 5302725, 9334867, 8078753, 11014223, 2694810, and 11299398. Shown in aqua in the lower panel are the 12 stars that MTF17 identified as exhibiting clear activity cycles over 4 years of observations (their Figure 5 ). Four of these are very clear-cut cases of spot dominance (those with the most negative r coefficients and slopes), two are unclear (those above the lower dotted line), and the rest are distributed along a continuum of spot dominance. Figure 11. Upper panel: five illustrative sets of brightness vs. variability. They have been scaled and shifted to fit on the plot. Color and symbol size are related to slope (see text). Lower panel: slopes of the linear fits to the brightness vs. variability points for all stars; the ordinate is the Pearson-r coefficient for each star. The colors indicate the designations from MTF17 based on the sign of their correlation slopes (see text). Green points are stars with clear brightness cycles in MTF17. MTF17 divided the sample into three cases: anti-correlated (red), correlated (blue), and not clear (black); the colors in Figure 11 correspond to their classification. As explained in Section 2.1 , the “anti-correlated” label is used for stars that become brighter as they become more variable; they are generally interpreted to be faculae-dominated. The points for such stars have a positive slope, so they should lie in the upper right quadrant of the lower panel of Figure 11 , which shows the results for the full sample. Similarly, the blue (correlated) stars get darker as they get more variable (spot-dominated) and are in the lower left quadrant. Most stars from the sample (320 in MTF17; 322 for us) are in this category. In most cases (310 stars), their sense agrees with ours, although there are a few red points in the lower left quadrant, and stars that they did not classify in either category (black) are found in both quadrants for us. The most likely source of discrepancies
Paper: 2018A&A...609A..29G (fulltext)
with the update of Goriely (1999) . The reactions and -decays are also introduced when relevant, in particular for the Bi and Po isotopes. 3. Nucleosynthesis process during core He-burning Fig. 2 Upper panel : Kippenhahn diagram of our 2 during the first dredge-up and core He burning phases; the green and blue dotted lines delineate the region of maximum nuclear energy production by H- and He-burning, respectively. Middle panel : evolution of the central neutron density , central C mass fraction () and neutron exposure . Lower panel : evolution of the central temperature and density . When the star leaves the main sequence and the central temperature reaches K (for a density ), the burning of the CNO by-product C via produces a first burst of neutrons (Fig. 2 ). This occurs during the first dredge-up of our 2 model star and leads to a maximum neutron density of and a time-integrated neutron exposure of , where is the most probable relative neutron-nucleus velocity at the temperature . The main neutron emission lasts until is depleted at the center and then proceeds at a much lower rate in the contracting shells surrounding the -free core. Eventually helium ignites off-center and after a series of shell flashes, convection reaches the center at an age yr. The temperature in the core reaches K and can trigger the production of a small amount of neutrons by the activation of the reaction. The neutron exposure increases up to but remains too small to generate an efficient -process nucleosynthesis. Instead, it is responsible for the re-arrangements of the abundances of specific heavy nuclei. As shown in Fig. 3 , elements such as K, Cd, Xe, Gd, Dy, Yb or Hg can be significantly produced with abundances increased by more than 0.4 dex with respect to their initial value and others partially destroyed, like Ag, Sm or Eu. At the temperatures encountered in the stellar center during core He burning, we notice that Te, Lu, Ta and Os become -unstable and fully decay ( Takahashi Yokoi 1987 ). When helium is exhausted in the core, the chemical profiles in the central 0.4–0.5 are frozen and the abundances of these heavy nuclei will be part of the future white dwarf. We note that the distributions of overabundances of elements heavier than typically show very small differences between our 2 and 3 stellar models, almost independently of the metallicity. Fig. 3 Final overproduction factors [X/Fe] of all stable nuclei with (full squares) at the center of our 2 model star at the end of core helium burning. The s-only nuclei are shown by red diamonds connected by a solid line. The vertical line depicts the total decay of the s-only nucleus Te. 4. <italic>S</italic>-process nucleosynthesis during the AGB phase of a 2 <italic>M</italic><inline-formula specific-use="simple-math"><sub>⊙</sub></inline-formula> [Fe/H] = –0.5 model star Mixing of the nuclear species beyond the convective boundary in the LM and DM models essentially depends on three parameters: , and .
Paper: 2021A&A...650A..17F (fulltext)
of the solar wind at 1au is a few km s −1 ( Němeček et al. 2020 ), which is small compared to the average radial wind speed of 400–600 km s −1 . Despite this, attempts to constrain the solar wind AM flux have been made in the past (e.g. Lazarus Goldstein 1971 ; Pizzo et al. 1983 ; Marsch Richter 1984 ; Finley et al. 2019 ). Though each of these studies suffered from uncertainty, either due to the precision of the spacecraft pointing or the ability of the plasma instruments to detect the small tangential flows. These uncertainties are lessened with Parker Solar Probe (PSP) which samples the solar wind within 0.25 au, where the signal to noise on the tangential wind speed measurements should be increased and the influence of wind-stream interactions reduced. However, during its first few encounters, PSP detected tangential solar wind speeds of up to ~ 50 km s −1 ( Kasper et al. 2019 ). These speeds are far greater than expected from magnetohydrodynamic (MHD) modelling (e.g. Réville et al. 2020 ) and, if they were prevalent throughout the corona, would be incompatible with the AM-loss rate implied by the slope of the rotation-age relation. Finley et al. (2020) analyse the data from the first two orbits of PSP and show these flows to be localised, existing along with wind streams that have large negative tangential speeds. Therefore, the average equatorial solar wind AM flux measured by PSP is similar to observationsfrom previous spacecraft. At present, the mechanism(s) that generate these large tangential solar wind speedshave yet to be determined. One possibility is the motion or circulation of open magnetic flux at the base of the wind ( Crooker et al. 2010 ; Fisk Kasper 2020 ; Macneil et al. 2020 ), though this is also not well understood. During PSP’s first two perihelion passes (E01 and E02, respectively), neither the Solar Probe Cup (SPC, Case et al. 2020 ) or the Solar Probe ANalysers (SPAN; Whittlesey et al. 2020 ; Livi et al. 2020 ) were able to confidently measure the properties of the alpha particles in the solar wind. Previously, Pizzo et al. (1983) used Helios observations to show that the alpha particles contain a large negative contribution to the total AM flux (see also Marsch Richter 1984 ). Though, due to an error in spacecraft pointing, the data from Helios required a significant correction, and more recent measurements of the alpha particle AM flux using the Wind spacecraft at 1au instead suggest a much smaller (but still negative) contribution ( Finley et al. 2019 ). As the alpha particles carry roughly 15− 40% of the linear momentum in the solar wind, it is vital that their contribution to the total solar wind AM flux is better understood in order to further constrain the slope of the rotation-age relation for Sun-like stars. In this work, we analyse more recent SPAN-Ion observations from the third and fourth encounters of PSP (hereafter E03 and E04), which we use to determine the mechanical AM flux of both the solar wind