adsabs/WIESP2022-NER huggingface dataset:
Paper: 2015ApJ...802...62D (acknowledgments)
We thank an anonymous referee for a very helpful report that enabled us to significantly improve the manuscript. We thank the NASA AISRP for providing financial assistance for the development of the PINTofALE package and the CHIANTI project for making publicly available the results of their substantial effort in assembling atomic data useful for coronal plasma analysis. J.J.D. and V.L.K. were supported by NASA contract NAS8-03060 to the Chandra X-ray Center during the course of this research. DG was supported by Chandra grants GO1-2021, GO2-3010X, and AR4-5002X.
Paper: 2018ApJ...853...69C (fulltext)
and abundance of galaxies in the Spitzer Matching survey of the UltraVISTA ultra-deep Stripes (SMUVS) to understand the connection between galaxies and their host dark matter halos. SMUVS provides a unique combination of large area (∼0.66 square degrees of the COSMOS field) and unparalleled depth at 3.6 and 4.5 μ m. This, combined with the latest ultra-deep UltraVISTA data and other ancillary data in the COSMOS field, allows for precise photometric estimates of redshift and stellar mass and robust statistics from such a large catalog ( objects). As a result, our analysis is performed over an unprecedented redshift range () for volume-limited galaxy samples selected by stellar mass. To interpret the observed clustering, we use a well-established five-parameter halo model, although we fix the values of three of these parameters. Novel to our approach is the inclusion of the photometric redshift dispersion, , as a free parameter in the fitting procedure. This parameter is designed to encapsulate the effect of dispersion in the photometric redshifts used in our analysis and allows us to fully exploit the clustering information in our photometric galaxy catalog. This parameter is discussed in more detail in Appendix A . From our best-fit halo models we can derive a number of interesting quantities. The characteristic halo masses, and , are well constrained by our data and exhibit fairly tight linear relationships (in logarithmic space) with stellar mass. This is similar to what has been found in earlier studies with similar analyses (e.g., McCracken et al. 2015 ). The satellite fraction, i.e., the fraction of galaxies in a sample that do not sit at the center of the potential well of their halo, generally increases with cosmic time. This is in broad agreement with other studies in the literature and in line with expectations of hierarchical structure formation. The comoving correlation length, r 0 , [defined such that ], and large-scale effective galaxy bias, , show strong monotonic trends with increasing redshift for a fixed stellar mass over the entire redshift range probed. Projecting ΛCDM halo mass accretion histories computed according to Appendix A of McBride et al. ( 2009 ) through our measurements of the large-scale bias, and combining this with the evolution of stellar mass content based on the cumulative number density (Behroozi et al. 2013a ), we propose an evolutionary path for the dark matter and stellar mass content of galaxies for . This shows a favorable agreement with our independent determinations of the relationship between galaxy bias and stellar mass at these redshifts and appears to be an interesting method to investigate the coevolution of these quantities in the future. Finally, we investigate the SHMR. This indicates the efficiency with which baryons have been converted into stars over the lifetime of a dark matter halo. We fit the SHMR parametrization of Behroozi et al. ( 2013b ) for our two lowest redshift bins ( and ).
Paper: 2019MNRAS.482.3803P (fulltext)
the SDSS data from at least one star-forming region in each galaxy in the sample suggest that the sample galaxies are undergoing young massive star formation phase having a significant population of WR stars. Figure 8. The blue WR bump detected in the spatially resolved star-forming regions of SBS 1222+614. The dashed-dot line represents the continuum fit. The solid line represents a Gaussian fit to the blue bump. Figure 9. The detection of the red WR bump in the spatially resolved star-forming region of SBS 1222+614. The dashed-dot line represents the continuum fit. The solid line represents a Gaussian fit to the red bump. 3.3 Ionization mechanism The detection of high ionization emission lines of He ii λ4686, [O iii ] λ4363, [O iii ] λ4959, 5007, and [Ar iii ] λ7135 and low ionization emission lines of [O ii ] λ3726, [N ii ] λ6584 and [S ii ] λ6717, 6731 etc., in the galaxy spectrum can be an indication of mix of many excitation sources such as photoionization mainly from massive ionizing stars and shocks generated by WR stars, supernovae explosions, and energetic active galactic nuclei-related mechanisms as discussed by Garnett et al. ( 1991 ), Garnett, Kennicutt Bresolin ( 2004 ), Guseva et al. (2000 ), and López-Sánchez Esteban ( 2010a ). Therefore, it is necessary to identify the dominant excitation mechanism in these galaxies. The nature of the dominant ionizing source can be inferred using the quantitative classification scheme proposed by Baldwin, Phillips Terlevich ( 1981 ) using combinations of line ratios. Subsequently, other similar diagnostic schemes were proposed by Dopita et al. ( 2000 ) and Kewley et al. ( 2001 ). All these diagnostic schemes use the emission line ratios of [O iii ] λ5007/H β, [N ii ] λ6583/H α and [S ii ] λ6717, 6731/H α. The optical emission line ratios for all the spatially resolved ionized regions in our galaxy sample are listed in Table 3 and the locations of all the star-forming regions are presented in Fig. 10 , which shows typical plots for [O iii ] λ5007/Hβ versus [N ii ] λ6583/Hα (top panel) and [O iii ] λ5007/H β versus [S ii ] λ6717, 6731/H α (bottom panel). Here, we used the latest Kewley et al. ( 2001 ) diagnostic criteria to identify source of excitation mechanism. In these figures, it can be seen that the primary and dominating source of ionization is photoionization in all the cases. Figure 10. A comparison of the observed flux line ratios obtained from the star-forming regions in our sample of dwarf galaxies with the diagnostic diagrams proposed by Kewley et al. ( 2001 ). 3.4 Physical conditions of the ionized gas The electron temperature ( T e ) and density ( n e ) of the ionized gas in the H ii regions are estimated here. Since the faint auroral [O iii ] λ4363 emission line along with the [O iii ] λ4959, 5007 emission lines were detected in many cases, the electron temperature could be estimated directly using their relative line intensities. A two-zone approximation was assumed to
Paper: 2020ApJ...898...91W (fulltext)
work is different from traditional KAW theory of , but is close to the theoretical prediction of from Boldyrev Perez ( 2012 ) and Boldyrev Loureiro ( 2019 ) and the simulation from Cerri et al. ( 2019 ). It also corresponds to α = 1.13 in the framework of the model by Landi et al. ( 2019 ). Hence a modified CB premise may be needed to understand the kinetic turbulence in magnetosheath. We note that, for most of the events, the weak and scale-independent anisotropy in the perpendicular plane appears to agree with recent PIC simulation of KAW turbulence (D. Grošelj 2020, private communication) but is inconsistent with results of Boldyrev Loureiro ( 2019 ), which predicts a much stronger anisotropy and a steeper scaling of or L ⊥ ∝ , depending on different current sheet configurations for the tearing instability. Capturing the active signatures of current sheet disruption/reconnection (i.e., Mallet et al. 2017 , Loureiro Boldyrev 2017 ) from in situ observation is a challenging task, but will contribute to understanding its effects on the anisotropy. The spectral anisotropy of kinetic plasma turbulence is believed to be associated with dispersion and intermittency effects (Zhao et al. 2016 ; Landi et al. 2019 ). For example, the anisotropic scalings are different below and above ion cyclotron frequency and also differs for sheet and tube-like turbulence (Zhao et al. 2016 ). To illustrate the possible connection between intermittency and spectral anisotropy, we specifically compare the results from two events. Figure 5 shows the excess Kurtosis and anisotropy relation for the total magnetic field. The excess Kurtosis is defined as , where S 4 ( l ) is the fourth-order structure function. In both panels of Figure 5 , the solid lines represent the results from event 1, which is recorded on 2017 October 4 and is used as our example of spectral anisotropy in Section 3 , while the dashed–dotted lines represent the results from event 2, which is recorded on 2017 December 24. As plotted in three directions l ⊥ , L ⊥ , l ∥ , the value of K ( l ) is around zero at large scales, meaning the roughly Gaussian distribution of the fluctuations. Toward small scales, K ( l ) displays an increase tendency before it drops down. Such non-Gaussian statistics ( K ( l ) > 0) confirms the presence of intermittency in the magnetosheath, while the scale-dependent profile of Kurtosis is similar with solar wind observations (He et al. 2019 ). For the kinetic-scale parallel–perpendicular anisotropy, we find that it can be considerably affected by the intermittency. As shown in Figure 5 , the stronger the intermittency (see the larger Kurtosis of the solid lines in the left panel), the stronger the anisotropy level (see the larger vertical distances between the solid blue curve and the gray line in the right panel). This phenomenon is consistent with previously observations at large scales, which emphasize the key role of intermittency in generating the spectral anisotropy
Paper: 2017MNRAS.466.1213K (fulltext)
star formation rate per unit area , if we have one supernova per mass M SN of stars formed, then the supernova rate per unit area is . The momentum and energy injected per supernova are p SN and E SN , respectively, giving momentum and energy injection rates per unit area and . The energy injection rate into turbulent motions is (equation 17 ) and the ratio of this to the total supernova energy budget is where the numerical evaluation is for our canonical values p SN 310 5 M kms 1 and E SN 10 51 erg. Thus, for the values of found in our simulation, only a small portion of the supernova energy budget is consumed by driving turbulence. Similarly, the wind kinetic luminosity per unit area is of the order of and the ratio of this to the supernova energy injection rate is where the numerical evaluation is for M SN 100M and v esc 300kms 1 , and is the mass loading factor, which for our simulations is 1. Thus we conclude that neither launching the wind nor driving the turbulence consumes an appreciable fraction of the total supernova energy available. Instead, the energy released by supernovae must either be lost to radiation, or must go into a hot wind that carries it out. Unfortunately, our model does not allow computation of the partition between these two forms of energy loss, and the simulations that have been published to date are not helpful in addressing this question answering it correctly requires simulating with enough resolution to resolve the SedovTaylor phase of supernova remnant expansion, without using any artificial methods to lower the density in the vicinity of the supernovae. As we discuss in Section 4.6 , no published simulations meet these criteria. Observations of superbubbles away from galactic centres suggest that radiation cannot be the primary loss mechanism (Rosen etal. 2014 ), but it is unclear whether we can generalize these conclusions to the very different environment of the CMZ. If radiative losses do not dominate, then the hot wind must carry an energy flux of order , and given our result that 50percent of the incoming gas is converted to stars, this implies an energy flux ergs 1 . This material will mostly be launched from the star-forming ring at 100pc. The non-thermal particle energy injection rate should be 10percent of the hot gas energy budget, implying ergs 1 . Because non-thermal particle acceleration happens even if the hot gas does not vent, and because the non-thermal particles have long mean-free paths and thus should be able to escape the disc, this estimate should hold even if the primary loss mechanism is radiation rather than a hot wind. Cosmic ray escape is confirmed observationally (Crocker 2012 , and references therein): the diffuse non-thermal emission from the CMZ in radio continuum and -ray bands implies that only a small fraction (10percent) of the power injected into non-thermal particles accelerated in the CMZ is lost in situ radiatively; most of this power is carried off by escaping cosmic
Paper: 2015MNRAS.452.1914G (acknowledgments)
HGM is grateful for support from the UK Science and Technology Facilities Council through the grant ST/I001204/1. LV is supported by the European Research Council under the European Community's Seventh Framework Programme grant FP7-IDEAS-Phys.LSS and acknowledges Mineco grant FPA2011-29678- C02-02. JN is supported in part by ERC grant FP7-IDEAS-Phys.LSS. AC is supported by the European Research Council under the European Community's Seventh Framework Programme grant FP7-IDEAS-Phys.LSS. WJP is grateful for support from the UK Science and Technology Facilities Research Council through the grant ST/I001204/1, and the European Research Council through the grant ‘Darksurvey’. Funding for SDSS-III has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation and the US Department of Energy Office of Science. The SDSS-III website is http://www.sdss3.org/ . SDSS-III is managed by the Astrophysical Research Consortium for the Participating Institutions of the SDSS-III Collaboration including the University of Arizona, the Brazilian Participation Group, Brookhaven National Laboratory, University of Cambridge, Carnegie Mellon University, University of Florida, the French Participation Group, the German Participation Group, Harvard University, the Instituto de Astrofisica de Canarias, the Michigan State/Notre Dame/JINA Participation Group, Johns Hopkins University, Lawrence Berkeley National Laboratory, Max Planck Institute for Astrophysics, Max Planck Institute for Extraterrestrial Physics, New Mexico State University, New York University, Ohio State University, Pennsylvania State University, University of Portsmouth, Princeton University, the Spanish Participation Group, University of Tokyo, University of Utah, Vanderbilt University, University of Virginia, University of Washington and Yale University. This research used resources of the National Energy Research Scientific Computing Center, which is supported by the Office of Science of the US Department of Energy under Contract No. DE-AC02-05CH11231. Results are based on observations obtained with Planck ( http://www.esa.int/Planck ), an ESA science mission with instruments and contributions directly funded by ESA Member States, NASA and Canada. Numerical computations were done on Hipatia ICC-UB BULLx High Performance Computing Cluster at the University of Barcelona.
Paper: 2016A&A...588A..56B (acknowledgments)
D.B. and T.R. dedicate this paper in gratitude to their long-term collaborator and friend Stanislav (Stan) Štefl (1955–2014), with whom this project had been planned like so many others before. The authors thank the BRITE operations staff for their untiring efforts to deliver data of the quality that enabled this investigation. They appreciate the permissions by Tahina Ramiaramanantsoa et al. to mention some early results of their ongoing work on ζ Pup . Coralie Neiner is thanked for the early communication of magnetic-field measurements with ESPaDOnS of Cen. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France. This research has made use of NASA’s Astrophysics Data System. ACC acknowledges support from CNPq (grant number 307076/2012-1) and FAPESP (grant number 2015/16592-0). GH thanks the Polish NCN for support (grant 2011/01/B/ST9/05448). A.F.J.M. is grateful for financial aid from NSERC (Canada) and FRQNT (Quebec). A.P.i. acknowledges support from the Polish NCN grant No. 2011/03/B/ST9/02667. A.P.o. acknowledges support through NCN grant No. 2013/11/N/ST6/03051. G.A.W. acknowledges Discovery Grant support from the Natural Sciences and Engineering Research Council (NSERC) of Canada. The authors from Poland acknowledge assistance by BRITE PMN grant 2011/01/M/ST9/05914.
Paper: 2017MNRAS.470..224T (fulltext)
of ionizing photons drops rapidly, and is therefore quite low when the supernovae start to create clear channels. At these early times, the amount of photons reaching the IGM before the supernova phase is low as well. Similarly, the lower panel of Fig. 9 shows the ionizing duty cycle as a function of halo mass, and , for four regularly log-spaced mass bins centred on log ( M ⋆ /M ⊙ ) = 7.75, 8.25, 8.75 and 9.25. This is in strong disagreement with the recent results of Kimm et al. ( 2017 ), who simulated the formation of mini-haloes and found that f esc can be very high even before the first supernova. However, their study focuses on much lower mass haloes, reaching M vir ≤ 10 8 M ⊙ at most. They also include a description of pop III stars and molecular hydrogen, which can lead to significant differences for haloes below the atomic cooling limit of M vir ∼ 10 8 M ⊙ . Figure 9. Ionizing duty cycle of the two most massive haloes, as a function of the galaxy (upper panel) and halo (lower panel) mass. The blue (red) dashed line shows the fraction of the time a galaxy of a given mass spends with higher than 10 50 (10 48 ) photons per second. On average, a more massive galaxy spends more time leaking LyC photons. On the other side of the mass spectrum, when M ⋆ > 10 6 M ⊙ , the duty cycle saturates, with and . This happens when the galaxy starts to reach a more stable regime, with feedback-regulated star formation, but where the star formation is rarely completely shut off. This means that the gas in the galaxy will usually not have the time to settle down in between two starburst events, leading on average to a more steady production of ionizing photons. 5 LOCAL ESCAPE OF IONIZING RADIATION In Section 4 , we have found that the escape fraction is mainly regulated by the interplay between supernovae, gas accretion and clustered star formation, and we have limited our analysis to the total escape fraction of each galaxy, computed at the virial radius. Because it is strongly linked with ISM-scale phenomena, we will now concentrate on the intrahalo processes leading to the escape of ionizing radiation. For this purpose, we will mainly focus on the most massive galaxy in our sample. 5.1 Does radiation escape from the emission clouds? Before reaching the IGM, the ionizing photons must travel through the complex distribution of gas inside and around the galaxy. Kimm Cen ( 2014 ), Ma et al. ( 2015 ) and Paardekooper et al. ( 2015 ) showed that when f esc is low, most of the photons are absorbed very locally, within 100 pc of their emission point, and that it is indeed crucial to resolve properly the ISM in order to study properly the escape of ionizing radiation. We find very similar results with our simulations, and we suggest that this is the reason behind the strong coupling between the supernova feedback and evolution of f esc . While the idea that clear channels for ionizing radiation from young stars are created by ionizing radiation itself is
Paper: 2019A&A...626A..33H (fulltext)
of the projection effect of looking from the side, this corresponds to x = 11 when viewed from directly above. The line profile over the EB, shown in the left panel , is located at [ x , y ] = [17, 11] Mm and is evident in the line wings. The line core, shown in the right panel of Fig. 4 , shows extensive fibrils, some have about the length of the computational box, 20 Mm, but no evidence of brightening directly above the site of the EB. The two loop systems that meet and are the source of the EB are visible in the center of the right panel of Fig. 4 . They stretch toward greater and lower y -values from their meeting point close to y = 11 Mm? When viewed from directly above, the EB region is bright in H α in the line core as well, but this is due to emission that is formed much higher in the atmosphere, at least 5 Mm above the photosphere: The cool canopy above the current sheet is heated by Ly α and Ly-continuum radiation from the hot plasma of the current sheet below, raising the temperature of the canopy plasma significantly, as is visible in the lower middle panel of Fig. 1 and in the left panel of Fig. 3 . This relatively high-temperature gas excites some H α emission in the region that is heated above the current sheet. This heating at large heights directly above the reconnection region is probably a result of the method we chose for the coronal back-radiation and raises the temperature more than what would occur on the real Sun. The heating does not affect the evolution or diagnostics of the EB or UV burst. Swedish 1-meter Solar Telescope observations (Ortiz et al., in prep.) show that the wing of the Ca II 854.2 nm line can be a very good proxy for the existence of EBs, with enhanced emission that is colocated and roughly cotemporal with bright H α wing enhancement. These observations in the blue wings of both the Ca II 854.2 nm and H α lines also show dark, presumably cool, surges that occasionally emanate from the site of EB emission. We find that the EB in this location shows brightening, starting at time t = 7500 s and lasting for at least 1200 s thereafter. This enhanced emission, 0.0735 nm in the blue wing of the Ca II 854.2 nm triplet line, is shown in Fig. 5 , which also shows emission in the line core and in the shape of the line profile. Similar as for H α , the cores of the Ca II H Kand the Ca II 854.2 nm triplet lines show the loops that emanate from the vicinity of the current sheet as narrow thin fibrils that stretch toward higher and lower y -values, approximately in the y -direction. The EB is visible as a narrow band of emission in the line wing, and it is also visible at the time of this image ( t = 8220 s) in the line core, although this emission is partially covered by overlying longer fibrils. Apparently rising from the EB location, about parallel with the positive y -axis, a small cool surge is seen to be accelerated along the field lines, or loops, that form the cool portion of the canopy. Signs of the