adsabs/WIESP2022-NER huggingface dataset:
Paper: 2019MNRAS.483.1537B (fulltext)
other hand (lower panel), the I b / I n ratio is only slightly dependent on C . 6 APPLICATION TO A REAL CASE: SN 1006 6.1 Why SN 1006 In the following, we shall apply the methods derived so far to a real case, refining them whenever required. The Balmer emission along the northwestern limb of SN 1006 represents for this an almost perfect case: observations with high spatial and/or spectral resolution, on which to test and use our models, are available; a shock velocity () high enough to limit the role of CE processes (this simplifies the models and makes them more reliable); structures near the limb that appear rather well ordered; no evidence of efficient CR acceleration, which justifies neglecting the effect of CRs in the analysis. With reference to the shock speed, Ghavamian et al. ( 2001 ) measured a spectral width of , from which for a low T e / T p ratio they derived a shock speed of . A lower shock velocity, , follows instead from the model of Morlino et al. ( 2013b ): for a more detailed and updated discussion, see Raymond et al. ( 2017 ). In the following analysis, we shall fix V sh to . With reference to CR acceleration, a low efficiency in the northwestern limb can be deduced from several results of observations: (1) the absence of non-thermal X-ray emission (Bamba et al. 2003 ; Winkler et al. 2014 ); (2) a non-detection of TeV (Acero et al. 2010 ) and GeV gamma-rays (Xing et al. 2016 ); moreover, from the analysis of Balmer emission, two additional pieces of information point towards a low CR acceleration efficiency: (3) the FWHM (full width at half-maximum) of the narrow line is (Sollerman et al. 2003 ), compatible with being produced by unperturbed ISM with T = 10 4 K, pointing toward the absence of a CR precursor able to heat the upstream plasma; and (4) an absence of evidence of any Hα precursor in front of the shock ( R + 07 ). 6.2 Some recent high-quality observations In this section we summarize some results of two relevant papers, Raymond et al. ( 2007 , hereafter R + 07 ) and Nikolić et al. ( 2013 , hereafter N + 13 ), in which Balmer emission from some areas along the northwestern limb of SN 1006 has been observed in great detail, and that contain a wealth of information, which will be used in the present work. The former paper is based on a very deep HST /ACS image covering most of the Hα emission from that limb, and the superb spatial resolution of that image (an FWHM of about 0.13 arcsec, corresponding to about , where d 2 is the distance of SN 1006 in units of ) allows one to probe scales shorter than the collisional mean-free paths, and therefore to resolve the physical structure of the filaments. The latter paper, instead, presents a map of a selected portion of the northwestern Hα limb of SN1006, roughly corresponding to ‘Regions 26–29’ (as dubbed by R + 07 ). The instrument used, VIMOS in the IFU mode on the Very Large Telescope (VLT), allows both good spatial resolution (with a pixel size 0.67 arcsec,
Paper: 2017ApJ...834..163A (fulltext)
2000 ). The long-term X-ray pulse profile and pulsed fraction evolution, such as that presented in Section 4.1 , could be evidence for the slow, plastic deformation of the crust needed to produce slower regions of superfluid (Woods et al. 2004 ; Thompson et al. 2002 ). Alternatively, regions of the outer core could be involved as a result of the strong magnetic fields. Magnetar strength fields provide a strong pinning of vortices to flux tubes in the outer core, which leads to a rotation lag between the normal and superfluid components, which, when relaxed, would lead to a spin-down event (Kantor Gusakov 2014 ). In typical rotation-powered pulsars, the expected internal nature of these spin-up glitches is argued to be indicated by the lack of observed radiative changes. This “radiatively quiet” condition does not hold for many glitches in magnetars. Many magnetar glitches are accompanied by drastic changes in their X-ray flux and pulse profiles (e.g., Kaspi et al. 2003 ; Dib et al. 2009 ). This could indicate that external forces are at play in radiatively loud glitches, and these external forces may lead to the net spin-down events reported in several magnetars, including the new event reported here. We have compiled a list of all published confirmed and candidate net spin-down glitches in Table 3 . In this list, we present the magnitude of the change in spin frequency, . For net spin-down glitches that are due to overrecovering spin-up glitches, the of the glitch is presented. For events where there is no sign of a recovery, we show an upper limit for the time during which the spin-down had to occur. We also note that 6 of the 7 reported net spin-down glitches are accompanied by radiative changes. 8 8 This increases to 6 of 6 if we ignore the disputed event; see Table 3 . This is a much higher fraction of radiative activity than seen in spin-up events, where only 5 of 22 timing anomalies were accompanied by radiative changes (Dib Kaspi 2014 ). Table 3 Reported Net Spin-down Glitches Notes. a Accompanied by decaying spin-up glitch with Δ ν = 10.8(4) × 10 −5 Hz. b Possibly accompanied by second glitch event. c Accompanied by decaying spin-up glitch with Δ ν = 2.0(4) × 10 −7 Hz. d Accompanied by decaying spin-up glitch with Δ ν = 5.1(5) × 10 −8 Hz. e Disputed events, see the notes below. f Not seen in analysis of same data by Dib Kaspi ( 2014 ). For the net spin-down glitches in which the net spin down is measured to be due to an overrecovery, we also report the recovery fraction, Q . Recovering glitches are often classified by their recovery fraction, defined as . In younger pulsars, when recoveries are observed, Q tends to be larger than in older pulsars (Wang et al. 2000 ). For the new glitch we report here, we measure . Glitches with have only been seen twice before: once following the magnetar-like outburst of the pulsar PSR J1846−0258 that is powered by a high magnetic field rotation, which had (Livingstone et al. 2010 ), and once during
Paper: 2018A&A...619A.162X (fulltext)
From the ∼56% remaining, 70% of them have a counterpart in either 1RXS or 2RXS catalogs. In this paper, we focus on and perform an extensive visual inspection and literature search of the other 30% of the systems. Of these objects, we describe in detail the analysis and characterization of a pilot sample of 13 candidates. We determined a spectroscopic redshift for each cluster candidate finding that most of them are local systems ( z ≲ 0.05), and are very extended (significant X-ray emission traced out to more than 10 arcmin radius). Moreover, all these objects have fluxes above the limit of the 1RXS and 2RXS catalogs. We argue that they might have been missed because of the source finder used in previous works, that is, the sliding-cell algorithm, which is not optimal for finding such low-flux, very extended objects. Our pilot sample also reveals that the surface-brightness distribution of the systems is flatter than the usual profile, that is β 2/3, and that the missed systems are located in the group regime given the determined masses from galaxy velocity dispersions or a L X − M scaling relation ( M 500 ≲ 10 14 M ⊙ ). In a subsequent paper, we will present and discuss in detail the remaining cluster candidates that we have found. Appendix A: Additional figure Fig. A.1. From left to right : RASS photon images ( left ), optical images, integrated count rates profiles, and spectroscopic redshift histograms ( right ) for the galaxy groups in our pilot sample. The optical images are from ATLAS (composed of g , r , i bands, group 2, 6, 10), SDSS (composed of g , r , i bands, group 1, 3, 7, 8, 9) or DSS ( r band, other groups). The images have a 1 × 1 deg 2 size, except for the ATLAS images, which are 0.30 × 0.30 deg 2 . The red cross shows the position of the group candidate. The magenta plus signs display the location of 1RXS or 2RXS sources. The cyan diamonds show the positions of galaxies used for redshift determination. The white asterisks show the positions of known optical galaxy groups/clusters obtained from NED. The green contours are obtained from the corresponding wavelet-filtered images, and the blue circle shows the extent value (i.e., the core radius) of the group. The integrated count rates profile is background subtracted (solid line). The dashed curves give the 1 σ statistical error of the count rates. The vertical and horizontal dot-dashed lines indicate the aperture radius and the count rates used for the primary photometric results. The dotted curve shows the β model of the cluster modelling used to extrapolate to r 500 (see Sect. 4.3.2 for further details). In the spectroscopic redshift histograms the redshift values are taken for all galaxies within a box of 1 × 1 deg 2 centered on the group position with z 0.10 from NED. The total number of galaxies used in each redshift determination is given by the second number displayed in the upper right corner of each panel (see also Table 3 ). The redshift distribution is fitted
Paper: 2016A&A...587A.121W (fulltext)
the point sources matched between the two shortest wavelength IRAC bands. To illustrate the locations of detected sources in our catalog, we present in Figs. 5 through 9 the 3.6 m image of each galaxy with the positions of detected sources overplotted, followed by an image of the galaxy without sources shown. For each galaxy, red points denote the location of point sources not likely to be associated with the galaxy, while green points demarcate the regions inside isophotes for the Kron radius for all galaxies but NGC 2366, where the isophote described in de Vaucouleurs et al. (1991) was used. Consequently, the green points are more probably actual sources associated with each galaxy. Figures 10 through 14 show color-magnitude diagrams (CMDs) of [3.6] versus [3.6][4.5] for each galaxy. The points follow the color scheme from the images showing each galaxy, with green points representing sources within the isophotes for with each galaxy, and red points sources in the field of each galaxy. For each magnitude bin, the average uncertainty was computed for both the [3.6] magnitude and [3.6][4.5] color, and is plotted as an ellipse on the left hand side of the plot. For reference, lines are drawn to show the completeness limit of mag and the color [3.6][4.5] = 0, where most stars lie in the plot. In terms of absolute magnitudes, the [3.6] limit corresponds to detections of all objects brighter than in every galaxy. One other interesting feature in the plot was explained by Khan et al. (2015c) . For both NGC 55 and NGC 253, the two galaxies with the most sources of those studied in this work, one can see a “break” at about the same magnitude as the completeness limit. This feature is representative of the two different populations in the photometry, one physically associated with the galaxy, and another simply sharing the same region of sky. Khan et al. (2015c) calls the second population “contaminants”. This population contains foreground stars near the [3.6][4.5] color of 0, some Milky Way Halo late-type giants and unresolved background galaxies in other regions of the CMD. Red supergiants (RSGs) have been empirically shown to exist in the region of negative color space ( Bonanos et al. 2009 , 2010 ). This is due to a suppression of the flux in the 4.5 m photometry owing to molecular bands, most prominently CO ( Verhoelst et al. 2009 ). However, spectroscopic investigations into the individual stellar components of Local Group dwarf irregular galaxies have discovered 15% of new RSGs at positive colors in CMDs for those galaxies ( Britavskiy et al. 2014 , 2015 ). Fig. 10 Colormagnitude diagram for sources in the field of NGC 55. All photometrically measured sources are marked with a red dot, while those residing within the isophotal limits of the galaxy are represented by green dots. Overplotted are lines denoting the limiting magnitude of 18.0 and a [3.6][4.5] color of 0. On the left are ellipses corresponding to the average uncertainty for
Paper: 2021MNRAS.500..558M (fulltext)
by the FWHM. Given the nearly symmetric double-peaked profile of A370p_z1 and the small peak separation, we restrict our analysis to relatively small N H i and v exp values because it is well known that high column densities and shell velocities would significantly broaden the line and erase the blue peak, respectively (Verhamme, Schaerer Maselli 2006 ). We perform a quantitative comparison between the observed line profile and the models using the χ 2 statistic. We find that models minimizing the reduced χ 2 preferentially select low N H i (), static geometries ( v exp = 0), low dust content, small b values ( b = 20) and relatively broad input lines (). We show the best-fitting model in Fig. 4 which corresponds to the following parameter set: . We can derive the LyC escape fraction from the best-fitting column density = 99 per cent, where σ 912 = 6.35 × 10 −18 is the H i photoionization cross-section at the Lyman limit. While several models could match the positions of peak emission within the errors, the model is the only one to reproduce the shallow central depression of the profile which is the important signature of a low H i opacity (Fig. 4 ). Searching a finer parameter grid is beyond the scope of this paper, but we note that even when adopting a higher column density cm −2 , the escape fraction remains very high (94 per cent). Figure 4. Comparison of RASCAS models, smoothed by the resolution of X-Shooter, of Lyman α transfer through non-expanding shells of homogeneous gas with the observed Lyman α profile. The best-fitting model (blue) has which correspond to a high LyC escape fractions . We also show examples of models with (dashed orange, dash–dotted green), which would imply . 3.2 Did A370p_z1 self-ionize its local H <sc>ii</sc> bubble? The detection of the blue peak in the Lyman α indicates A370p_z1 sits in a large ionized bubble, otherwise the damping wing of even a partially neutral IGM would have absorbed it. Given its high escape fraction, we now consider whether A370p_z1 could have self-ionized its local H ii bubble. The blue wing extends to λ = 9479.2 Å, from the line centre, corresponding to a physical distance r H ii > 0.26 ± 0.04 pMpc. This estimate neglects any velocity offset between the Lyman α absorption dip from which we have derived the redshift and the systemic redshift as defined by more reliable tracers such as nebular absorption lines. These velocities are however found to be (Gazagnes et al. 2020 ), which, in the worst case, would therefore require a larger ionized bubble for the blue wing to be transmitted (≲ 0.5 pMpc). We also neglect peculiar velocities of the galaxy with respect to the ionized bubble gas which would redshift(blueshift) the Lyman α photons and decrease(increase) the bubble size needed for the blue wing to escape. We now estimate the volume that could have been reionized by A370p_z1 by redshift z = 6.803 and whether its radiation is sufficient to reduce the opacity of the surrounding gas to
Paper: 2017ApJ...843...91L (fulltext)
is a mid-stage merger consisting of two face-on galaxies which appear to overlap and have a projected separation of ∼84 (6.3 kpc). Prominent rings/arms in each galaxy contain the bulk of the visible star clusters (Figure 19 ). The maximum A V adopted for this galaxy is 1.0 mag of visual extinction (see Figure 20 ; Miralles-Caballero et al. 2012 ). The observed and derived properties for the star clusters identified in this galaxy are given in Tables 15 and 16 , respectively. 2MASX J08370182-4954302 2MASX J08370182-4954302 is a mid-stage merger containing two nuclei separated by ∼066 (0.36 kpc). Surrounding the nuclei are multiple bright star clusters in a spiral ridge just northwest and west of the nuclei (Figure 21 ). The maximum A V adopted for this galaxy is 3.7 mag of visual extinction (see Figure 22 ; Rich et al. 2012 ). The observed and derived properties for the star clusters identified in this galaxy are given in Tables 17 and 18 , respectively. NGC 2623 Evans et al. ( 2008 ) discuss the detailed morphology of this galaxy at length. NGC 2623 is a late-stage merger with dust lanes running along its tidal tails into the nucleus. Several bright clusters are distributed throughout the bulge and in a “pie-wedge” concentration south of the nucleus (Figure 23 ). The maximum A V adopted for this galaxy is 1.9 mag of visual extinction (see Figure 24 ; Privon et al. 2013 ). The observed and derived properties for the star clusters identified in this galaxy are given in Tables 19 and 20 , respectively. UGC 04881 UGC 04881 is an early-stage merger containing two nuclei separated by ∼11′′ (9 kpc). Spiral dust lanes and strings of star clusters surround the NE nucleus. In the SW nucleus, a linear distribution of star clusters and a prominent dust lane are seen (Figure 25 ). The maximum A V adopted for this galaxy is 1.9 mag of visual extinction (see Figure 26 ; González-Martín et al. 2009 ). The observed and derived properties for the star clusters identified in this galaxy are given in Tables 21 and 22 , respectively. IC 2545 IC 2545 is a late-stage merger being viewed face on. Dust lanes and strings of star clusters extend from two unresolved nuclei separated by ∼08 (0.54 kpc) in the center of the galaxy. Multiple star clusters are also visible throughout the tidal tails (Figure 27 ). The maximum A V adopted for this galaxy is 4.0 mag of visual extinction (see Figure 28 ; van den Broek et al. 1991 ). The observed and derived properties for the star clusters identified in this galaxy are given in Tables 23 and 24 , respectively. NGC 3256 NGC 3256 is a late-stage merger containing a large number of star clusters along the inner (∼20′′, or 4 kpc) spiral structure of the nuclear region. The spiral dust lanes extending from the nucleus give this galaxy pockets of high and low extinction (Figure 29 ). The maximum A V adopted for this galaxy is 3.3 mag of visual extinction (see Figure 30 ; Rich et al. 2012 ). The observed and derived
Paper: 2016ApJ...821L..23Z (fulltext)
field is almost horizontal and the field lines are concave-up. These are also locations where plasma can accumulate and eventually form a prominence (Aulanier Demoulin 1998 ; Lionello et al. 2002 ; van Ballegooijen 2004 ; Dudík et al. 2008 , 2012 ; Mackay van Ballegooijen 2009 ; Gibson 2015 ). The orange magnetic field lines extend up to z = 0.01 above the magnetic dips. If we scale our active region to the size of a typical active region hosting quiescent prominences, i.e., about 300 Mm, then orange magnetic field lines extend up to km above the magnetic dips. This is comparable with the (chromospheric) pressure scale height, therefore resulting in a reasonable proxy of quiescent prominences as the ones discussed in Filippov Den ( 2001 ). A comparison between each simulation suggests that different photospheric motions may result in prominences with similar shapes, but different morphologies. The side views of “Run C” show that the prominence is fragmented in its central part, i.e., it has a region devoid of dips at low altitudes, and it is also asymmetric: one of its ends is higher than the other one. This is not the case for “Run D1” where the distribution of the dips seems to follow a more homogeneous circular shape. One difference between the two cases is the asymmetry of the flows that formed the flux rope: the flows in “Run C” have a higher degree of asymmetry compared with “Run D1.” The details of the formation process may therefore influence the exact morphology of the flux rope and that of the dips, with these differences becoming more difficult to detect in the proximity of the flux rope’s axis ( ZAG15 ). The right-hand side panels show the evolution of the system at three different times. The configuration in the top panels is stable, the one in the middle is unstable in the low diffusion regime, but stable in the high diffusion regime, while the configuration in the bottom panel is unstable also in the high diffusion regime. The top-two horizontal yellow lines indicate the critical decay index and critical height for the onset of the instability under “high” and “low” diffusion conditions (see Table 1 ), while the bottom lines in all three models indicate the height where the decay index is n = 1.1. For “Run C” (respectively, “D1,”“D2”) at the time of the onset of the eruption in “low” diffusion regime, the apex of the axis of the MFR, i.e., the location where the stability criteria should be evaluated, has reached a height where the decay index is n = 1.4 (respectively, ) while the highest magnetic dips reach a height where n = 1.1. Interestingly, this value n = 1.1 is independent of the particular formation process in the three cases presented here. This suggests that the apparent instability criterion as evaluated at the location of the prominence is less sensitive to the flux rope formation process compared to when it is evaluated at the axis of the flux rope. Similarly to the flux ropes discussed in ZAG15 , these MFRs are
Paper: 2016MNRAS.463.3339T (acknowledgments)
We thank Chris Hales, Anita Richards, Neal Jackson, Lee Whittaker and Joe Zuntz for useful discussions. We further thank Chris Hales, Eva Schinnerer and Vernesa Smolcic for providing us with the partially processed VLA data. MLB is an STFC Advanced/Halliday fellow. BT, IH and MLB are supported by an ERC Starting Grant (grant no. 280127 ).