adsabs/WIESP2022-NER huggingface dataset:
Paper: 2021MNRAS.505.4857P (fulltext)
cut through the current sheet in turbulence, the parameters depend on time due to the turbulent interaction of larger scale magnetic eddies, so wide variation in parameters across different current sheets and at different times is expected. 7 DISCUSSION AND CONCLUSIONS As a result of the weak collisionality of many space and astrophysical plasmas, several physical processes can locally contribute to the dissipation of energy and the heating of the plasma. During magnetic reconnection, magnetic energy is efficiently converted to directed plasma flows, thermal energy, and energetic particles (Burch et al. 2016 ; Torbert et al. 2018 ). During turbulence in a magnetized plasma, the cascade provides the global amount of energy needed for fluctuations to be dissipated at small scales (Marino et al. 2008 ). However, which dissipative mechanisms are dominant and under which conditions is still not sufficiently understood (Vaivads et al. 2016 ). Therefore, a multiplicity of dissipation surrogates have been adopted in the literature to identify potential sites of dissipation. Owing to the strongly non-linear dynamics and the importance of physics at kinetic scales in such systems, numerical simulations of the Vlasov–Maxwell system (or, including collisions, the Boltzmann–Maxwell system) are the decisive tool to address the long-standing issue of plasma heating and energy dissipation in magnetized plasmas, e.g. Parashar et al. ( 2009 ). For example, the ‘turbulence dissipation challenge’ (Parashar et al. 2015 ) motivated numerical work to compare different algorithms on the solution of similar problems to assess the nature of the dissipation in magnetized turbulence (e.g. Pezzi et al. 2017 ; Perrone et al. 2018 ; González et al. 2019 ). In the same spirit, we have conducted a survey of a number of dissipation surrogates with three different kinetic plasma codes – the fully kinetic particle-in-cell vpic , the fully kinetic Eulerian Vlasov–Maxwell Gkeyll , and the Eulerian Hybrid Vlasov–Maxwell codes – to perform numerical simulations of two important physical phenomena in plasma physics. The first class investigates reconnection in an isolated current sheet, and the second concerns plasma turbulence at kinetic scales. We have calculated and compared eight distinct dissipation proxies, delineated in Section 2 . For the sake of clarity, we have categorized them in terms of (i) energy-based parameters, whose definition describes energy transfer and conversion; and (ii) VDF-based parameters, that are directly related to kinetic signatures in the particle VDF. Energy-based parameters considered here are the power density by electromagnetic fields on charged particles (Zenitani et al. 2011 ), the pressure–strain interaction (Yang et al. 2017a ), and a local proxy of the turbulent energy transfer (Sorriso-Valvo et al. 2018a ). The VDF-based parameters are the local pressure agyrotropy (Swisdak 2016 ) and three measures describing how different a local
Paper: 2021ApJ...908L..44L (fulltext)
be generated at levels within the cloud deck, where the background continuum emission originates. It must be generated well above this layer, where the absorbing gas is cooler and therefore absorbs more efficiently than it emits. The narrow width of the absorption line also suggests that it was formed at pressures substantially less than those of the cloud top (70 km, ∼3000 Pa). Figure 4. Radiance mixing ratio Jacobians ( dI ν / dr mix ) as a function of layer pressure for the radiance streams at 21 degrees zenith angle, for the 266.94 GHz feature for SO 2 (left) or PH 3 (right). Absorption line strengths (lower panels in each case) are given in units of cm −1 /(molecule cm −2 ). The continuum originates from SO 2 at ∼6 × 10 4 Pa (∼54 km, within the cloud deck), while the line cores for either species do not originate in the clouds (48–70 km) but at over 400 Pa (over 80 km) in the mesosphere. In the right panels in both plots, the temperature structure is given as a black line, while the colored lines denote SO 2 (green) or PH 3 (purple) mixing ratios. On the right, the evenly mixed 20 ppb PH 3 profile is shown with a solid line and the photochemical PH 3 profile is shown with a dashed line. 3.3. ALMA Line Dilution While the nondetection of prominent SO 2 spectral features in the ALMA wideband data could indicate a low abundance, as argued by Greaves et al. ( 2020a ), the estimation of this abundance was done without correcting for line dilution as a result of the ALMA observing geometry (Greaves et al. 2020a ). Significant line dilution is likely, especially considering the global distribution of SO 2 in the Venus atmosphere, and the exclusion of the short baseline ALMA measurements. Greaves et al. ( 2020a ) estimated line dilution (filtering losses) of 60%–92% depending on position on the disk. To determine the disk-averaged line dilution for the SO 2 reference line search, we simulated observations of Venus using the ALMA configuration of Greaves et al. ( 2020a ) by imposing an appropriate resolution spectrum (0.00003 cm −1 , 1 MHz) of our Case D atmospheric model over a limb-darkened disk model. The Fourier Transform of this model was resampled to match the ALMA configuration and reimaged using the imaging routines of Greaves et al. ( 2020a ), as provided in their Supplementary Software 3. As shown in Figure 5 , line dilutions on the order of 95% at the line core are observed for the full disk. We observe similar dilutions when the spectrum is only imposed on one hemisphere (∼8′′ extent at the time of observation). This line dilution suggests that the SO 2 reference features produced by our best-fit SO 2 distribution (Case D) would be heavily suppressed by line dilution in the ALMA data, which would cause them to mimic smaller features below the ripple detection limit of −0.0006. Figure 5. Modeled line dilution for the ALMA observations: disk-averaged line/continuum ratios for our nominal Case D atmosphere model containing SO 2
Paper: 2017ApJ...849...40G (fulltext)
our best guess associates the sunquake location (taken to be the 5 mHz signal in Figure 1 ) to the connectivity domain indicated as Q 4 in Figure 12 . That connectivity domain corresponds to the domain 4 in Figure 13 , i.e., to the envelope bundle of field lines (in green in Figure 14 ). Given its structure, such a field line bundle has two interesting properties. First, since it is surrounding the flux rope in almost its entire length it is then directly impacted when the eruption occurs. Second, such wide-spread field lines between the erupting flux rope and the above layered structure are all connected on one side to a very small area in the photosphere (where the sunquake occurs), in this way, providing a sort of magnetic focusing. In light of this, we argue that the sunquake is a consequence of the energy release associated with the reconnection between the erupting flux rope and the envelope surrounding it. Furthermore, the envelope structure helps focus the released energy into a small area in the lower atmosphere ( Q 4 in Figure 12 ), which produces the sunquake pulse. The scenario offered by our analysis and the interpretation of the observations clarifies the dynamics of the two flares and supports the association of the sunquake to the larger, M-class flare, rather than to the first C-class flare. Moreover, it indicates how a sunquake may require a special magnetic field configuration to occur, which may help understand their erratic occurrence. 6. Conclusion In this work, we study the evolution of active region 11675 on 2013 February 17, in the hours leading to the production of two flares (a C-class flare and an M-class flare), the ejection of filament material and a sunquake. We investigate the configuration of the active region magnetic field through the construction of an NLFFF extrapolation just over one hour before the sunquake. The aim of the work is to shed light on the configuration of the coronal magnetic field in which the sunquake was produced. The sunquake occurs when active region 11675 is undergoing a period of rapid evolution during its emergence phase, and when there are strong motions of the photospheric magnetic field elements. In particular, two emerging bipoles collide and shear past each other. The magnetic field along the polarity inversion line that separates these bipoles is highly sheared and at low altitudes bald patch field lines are present. At higher altitudes, an HFT and a flux rope are present, which have been highly modified by the motions of the photospheric field in which these magnetic structures are rooted. From a time–distance analysis, we determine that the sunquake was initiated in the time period between 15:50:41 to 15:51:26 UT on 2013 February 17 in a region of positive polarity field. This finding indicates that the sunquake is initiated during the M-class flare, which is associated with the eruption of the flux rope and ejection of filament material. However, this timing determination
Paper: 2021MNRAS.504.2983T (fulltext)
more spheroidal, as in the Milky Way halo, rather than being represented by a disc-like structure. One would therefore expect intermediate-aged populations traced, for instance, by RC stars to be somewhere between those. This is exactly what we found. Subramanian Subramaniam ( 2012 ) also used RC stars but found a very small inclination, i = 0.5°. Their methodology however was to use axes ratios and to limit the sample to a small depth range, which may have inadvertently introduced a bias. The PAs we found are nearest to those determined by Subramanian Subramaniam ( 2015 ), who used Cepheids, and Rubele et al. ( 2015 ), who fitted entire CMDs and found a mean i = 39°. In conclusion, while the SMC does not exhibit an obvious disc component there is a clear trend for the system to become more inclined with respect to the plane of the sky the younger the (stellar) tracer is. Since the RC stars already show this progression this means that the dynamical events that led to this time-dependent morphology must have occurred on Gyr time-scales. Indeed, while most of the stellar mass in the SMC was generated >4 Gyr ago (Sabbi et al. 2009 ; Rezaeikh et al. 2014 ; Rubele et al. 2018 ) there is evidence in the SFH of the SMC of a strong interaction with the LMC to have occurred ∼2 Gyr ago (Harris Zaritsky 2004 ; Noël et al. 2009 ; Rezaeikh et al. 2014 ; Rubele et al. 2015 ). A more recent interaction ∼0.2 Gyr ago coincident with the closest approach to the Milky Way, also seen in the SFH (Harris Zaritsky 2004 ; Noël et al. 2009 ; Sabbi et al. 2009 ; Rubele et al. 2015 ), may have caused further distortion of the gaseous disc, from which the younger population traced by the Cepheids formed. 6.3 Distortions and other structures The outer regions of the SMC have been transformed by interactions leading to stripping and spatially stretched populations. Dias et al. ( 2016 ) described the SMC as being composed of three bodies: the eastern Wing, the central main body, and the western halo. We find the intermediate-age main body traced by RC stars to be structurally simple and symmetrical, consistent with a spheroidal component. The eastern Wing, on the other hand, is split into two distinct populations (both traced by intermediate-age RC stars), of which the brighter one clearly lies in front of the main body (Fig. 9 , bottom right). We have been able to comprehensively trace these two components across the VMC footprint, allowing us to relate them to what other observational and theoretical studies have found. We find a fourth structure, in the north (Fig. 9 , top left); this has previously been described as an extension to the Wing (e.g. Gardiner Hawkins 1991 ), but Sun et al. ( 2018 ) showed this area to contain a shell of young stars which is clearly separated from the general Wing area. The intermediate-age population of RC stars in the north need not be related to these young stars, though. We compare our results with previous analyses of Cepheids, RR Lyræ,
Paper: 2016A&A...585A.149W (fulltext)
G327.29-0.5 – show shallow blue absorption wings that are not reproduced by the fits because no outflow component is included in the model. Nevertheless, the fit methods allow us to reliably adjust the free-fall fraction even in these cases. 4.4. Simultaneous HCO<inline-formula specific-use="simple-math"><sup>+</sup></inline-formula> modeling Even more indications of an additional outflow component are seen in most HCO lines. The dotted fits to the HCO lines in Fig. 6 show the infall models with only an adjustment of the HCO abundance. In general, they cannot reproduce the observed profile signatures, that is, the broad emission, the deep self-absorption, and the relative intensities of the blue and red peaks. We therefore added outflow emission to the RATRAN models of HCO in the center plane of the envelope before the ray-tracing step, as was done in Mottram et al. (2013) . Examples of profiles that can be fit satisfyingly with this additional component are G34.26+0.1 and G35.20-0.7, as shown in Fig. 6 . The outflow components have typical widths of 10 km s and are in some cases slightly shifted from the systemic velocities to account for the asymmetry of the outflow lobes. For several clumps, the HCO abundances are relatively small but consistent with typical depletion factors found in Giannetti et al. (2014) . Still, several HCO profiles cannot be reproduced with the simplicity of the spherical models, pointing to inherent difficulties in the interpretation of HCO profiles as an infall indicator. Fig. 6 Comparison of APEX HCO (4–3) observations with the SOFIA ammonia absorption in sources showing redshifted absorption. We show the RATRAN modeling result of HCO (4–3) in red using variable abundances with and without (dotted) additional outflow components. 5. Discussion and conclusions Adding the results from Wyrowski et al. (2012) , eight out of eleven massive clumps have been found with redshifted absorption that is indicative of infall motions. This fraction of 72% is substantially higher than found in past searches for the blue-skewed profile signature, see, for example, Fuller et al. (2005) , Wu et al. (2007) , López-Sepulcre et al. (2010) , Reiter et al. (2011) and Rygl et al. (2013) . From our comparison of ammonia absorption with line profiles of dense gas tracers lines, we conclude that outflows are the most likely reason for hiding infall motions in the clump envelopes. This effect is less prominent in the THz ammonia spectra. While we see the red lobe due to the high excitation even in emission in the extreme outflow source G5.89-0.4, other sources are, if at all, only affected by blueshifted absorption from the outflow in front of the continuum. These absorption wings are similar to the blue emission wings in the CS (7–6) spectra, but do not affect the measurements of the main redshifted absorption from the envelope. The excitation in most outflows is too low to excite the high critical density THz ammonia line (), but are high
Paper: 2019MNRAS.490.3440S (fulltext)
are very important to understand the ICME’s morphology (Jones et al. 2002 ), propagation and acceleration of energetic particles (Reames 1999 ; Sanderson et al. 2000 ; Manchester et al. 2005 ), and cosmic ray modulation, i.e. Forbush decrease (Raghav et al. 2014 ; Bhaskar, Subramanian Vichare 2016b ; Bhaskar et al. 2016a ; Raghav et al. 2017 ; Shaikh, Raghav Bhaskar 2017 ; Shaikh et al. 2018 ). In fact, the role of these fine structures during geomagnetic storms has also been studied independently (Kataoka et al. 2015 ; Palmerio et al. 2016 ; Raghav et al. 2018 ). Recently, we have reported a unique feature of the ICME sheath, i.e. the coexistence of PMS and Alfvénic fluctuations. Here, we study their influence on the topology of a geomagnetic storm profile in detail. 2 OBSERVATION The studied ICME crossed the Wind spacecraft on 2000 October 12. The temporal variation of in situ plasma parameters during the ICME transit is shown in Fig. 1 . We have used 1 min resolution in situ interplanetary data from the OMNI database available at https://cdaweb.sci.gsfc.nasa.gov/index.html/ . The ICME substructures, i.e. shock–sheath and MC, are shown in different colour bands (see Fig. 1 ). The onset of the shock front is identified as sudden sharp enhancement in IMF B mag , solar wind speed ( V p ), proton temperature ( T ) and proton density ( N p ) (see the first vertical black dashed line). The shock front is followed by high N p , T, V p , and plasma beta (β) and large fluctuations in the IMF vector (i.e. B x, y, z ) region, which corresponds to the ICME shock–sheath (Shaikh et al. 2017 ; Kilpua et al. 2017a ). The estimated thickness of the ICME shock–sheath is about 0.21 au, which is calculated by multiplying the average solar wind velocity (∼450 km s −1 ) by total time duration (∼19.5 h). Information about the shock at 1 au observed by the Wind spacecraft is available at https://www.cfa.harvard.edu/shocks/wi_data/00179/wi_00179.html . Their analysis suggests that it is a fast forward (FF) shock with the following average properties: the shock normal , and the angle between the shock normal and the magnetic field θ Bn = 71.5°. This implies that the shock is a quasi-perpendicular shock. After the passage of the sheath, we observe a decrease in fluctuations in B mag and B comp , a slow variation in θ and φ, and a slow steady trend in V p and low β, indicating the transit of MC (Zurbuchen Richardson 2006 ). The catalogue available at http://space.ustc.edu.cn/dreams/wind_icmes/ corroborates the identified boundary of MC. Figure 1. ICME crossed the Earth during 2000 October 12–14. The figure has seven sub-panels. The first and second panels show the temporal variation of IMF B mag and B vec (i.e. B x , B y and B z ). The third panel represents the azimuth (φ) and elevation (θ) angles of the IMF vector. The fourth panel expresses the temporal variation of the geomagnetic index, i.e. SYM-H and the A L index. The fifth and sixth panels show solar wind
Paper: 2021MNRAS.505.2186D (fulltext)
of the cluster will change with the cluster properties. For a massive cluster or for a compact cluster the binding energy will be higher and hence the formation of the SMS would be more favourable. On the 5 Myr time-scale explored in our simulations, supernova feedback is also expected to become relevant. With the typical energy of 10 51 erg for a core-collapse supernova, it is clear that one such event will expel the gas and terminate the accretion, if it has not stopped already (either due to gas expulsion by stellar winds or as the accretion process may have depleted the gas). In future work, it will be important to study detailed gas dynamics where the kinetic energy deposition of winds as well as the supernova feedback is taken into account. 5 SUMMARY AND DISCUSSION In this work, we explored the effect of mass-loss due to stellar winds on the final mass of the SMSs, which could be formed via runaway stellar collisions and gas accretion inside NSCs. We find that an SMS of mass could be formed even in a high-metallicity environment for high accretion rates of . For an accretion rate of , the final mass of the SMS for . Whereas for solar metallicity, no SMS can be formed for f Edd = 0.9 and SMSs of masses can be formed for f Edd = 0.7 and 0.5, respectively. For the case of Eddington accretion it will not be possible to form an SMS in the metallicity regime . Finally, for the Bondi–Hoyle accretion scenario, we find that the formation of a SMS will not be possible in the high-metallicity regime of . The interaction of the stellar wind and the gas inside the cluster might play an important role in the evolution of the SMS. The winds from the SMSs have high velocities ∼10 3 km s −1 (Muijres et al. 2012 ), which might exceed the escape velocity from the centre of our modelled star cluster. The SMS in the cluster is close to the centre which results in a high collision rate near the centre due to a shorter relaxation time in the core and an increased collisional cross-section. If the SMS is displaced by collisions, it rapidly sinks back close to the centre via dynamical friction where it may eventually decouple from the remainder of the cluster. This is also known as the Spitzer instability (Spitzer 1969 ). Interestingly, Krause et al. ( 2016 ) have found that for a Salpeter -type mass function the stellar wind cannot remove the gas inside the cluster. Hence, we do not expect the stellar wind to remove gas from the cluster. One of the main caveats of this work is the neglect of the kinetic energy deposition through stellar winds, which could contribute significantly to expel the gas. The latter is likely to create a regime where the growth of a massive object is still inhibited, even though the mass-loss itself from the winds is no longer significant. Below a critical metallicity in the range (, this effect is no longer expected to be relevant; however, supernova feedback may lead to the expulsion of the gas. Another relevant caveat is the
Paper: 2021ApJ...916...89R (fulltext)
astrometry between the initial image and each of the later sets using Astroalign . To uncover or place limits on transient emission on these timescales, we perform digital image subtraction using the HOTPANTS software package (Becker 2015 ). If an optical or NIR source is detected within the GRB localization regions (Evans et al. 2009 ; Lien et al. 2016 ), we perform aperture photometry on the residual images to determine the source brightness using the IRAF/ phot package. We determine zero-points by using point sources in common with the SDSS Data Release 12 (Alam et al. 2015 ), Pan-STARRS DR1 and DR2 (Chambers et al. 2016 ), and 2MASS (Skrutskie et al. 2006 ) or USNO-B1 (Monet et al. 2003 ) catalogs. If there are no significant residuals at or near the localization region, we obtain an upper limit from sources detected at 3 σ significance in the field. As we obtain the limit from the initial image, our values do not include additional noise from the subtraction. In total, we present 52 unpublished upper limits and three unpublished detections of 23 bursts, 10 with known redshifts (Table A1 ). 2.3. Observations from GCNs We search the GCNs of all short GRBs detected between 2013 and 2020 for constraining upper limits to ensure that our catalog covers all Swift short GRBs to date (2005–2013 is covered by the literature sample in Section 2.4 ). Considering that the range of peak kilonova luminosities based on models is ν L ν ≈ (1−5) × 10 42 erg s −1 (∼10 times the peak luminosity of AT 2017gfo; see Figure 1 ), we only include data that could feasibly fall into the range of kilonova luminosities, or fainter than a given depth, corresponding to griz ≳ 22.5 mag and JHK ≳ 21.0 mag (AB system) at z ≈ 0.25. We choose z ≈ 0.25, as 80% of known redshifts in our sample fall above this cut. Observations significantly brighter than this are not included, as they are likely dominated by afterglow emission. In total, we collect deep optical and NIR limits of 18 bursts, seven with confirmed redshifts, from the GCNs (Table A1 ). Figure 1. Light curves of AT 2017gfo and four diverse kilonova models considered in this work in the griJ bands. Light curves are shown in ν L ν (left y -axis) and absolute magnitude (right y -axis) units, and rest-frame time in days. In each panel, we show the interpolated, smoothed AT 2017gfo data set (solid dark-blue line; Villar et al. 2017 ), an NSBH kilonova model (blue dotted line; Kawaguchi et al. 2020b ), an infinite-lifetime NS remnant kilonova model (red dashed line; Kasen et al. 2015 ), an extreme optimistic magnetar-boosted kilonova model (turquoise solid line; Fong et al. 2021 ), and a model for an optimistic scenario motivated by GW190425 (magenta dashed–dotted lines; Barbieri et al. 2020 ). Black and open triangles denote upper limits of short GRBs with known and unknown redshifts, respectively, in the respective rest-frame bands at δt rest = 0.1–10 days. The most constraining limits of bursts with known redshifts