adsabs/WIESP2022-NER huggingface dataset:
Paper: 2017ApJ...834...42J (fulltext)
show a decrease in the distance between conjugate footpoints (Figure 10 ). The converging motion of flare ribbons is confirmed by some recent observations (Ji et al. 2006 , 2007 ; Liu et al. 2008 ; Joshi et al. 2009 ) and, thus, it has been identified as an important phenomenon occurring during the rise phase of flares. In particular, Liu et al. ( 2008 ) noted that during SOL2002-04-30, the two conjugate HXR footpoints first move toward and then away from each other, mainly parallel and perpendicular to the magnetic inversion line, respectively. Furthermore, the transition between these two phases of footpoint motions coincides with the reversal in direction of the motion of the looptop source. The explanation of ribbon motions of this kind is still beyond the scope of the standard flare model. After a prolonged rise phase, the flare exhibited the features of a “standard” large eruptive flare as described by the standard flare model. Here the main observational characteristic of the standard flare is recognized in the form of an increase in the separation of parallel flare ribbons while they move perpendicular to the PIL (Figure 11 ). Following the eruption of the flux rope, we observed the development of an arcade of coronal loops over the source region that was previously occupied by the coronal sigmoid. This phenomenon, called sigmoid-to-arcade transformation, is well established and theoretically studied in relation to the eruptive dynamics of CMEs (Gibson et al. 2002 , 2004 ). In the radio spectrum, the violent eruption of the flux ropes is followed by the occurrence of a broad, intense type II radio burst, implying the propagation of a shock wave associated with the passage of the CME (see, e.g., Cho et al. 2005 ). We, therefore, conclude that the standard flare model is well applicable to the late phase of this flare. We believe that, during this phase, the main driver of the energy release process is the large-scale magnetic reconnection driven by the erupting FR. 5. CONCLUSIONS Sigmoid-to-arcade development is considered to be an important aspect of solar eruptive phenomena. Although the association of sigmoids with eruptions is well recognized, our knowledge about the formation of these structures and their subsequent activation is still limited. In this paper, we provide a comprehensive multi-wavelength investigation of EUV, UV, H α , X-ray, radio, and magnetic measurements to probe the evolution of a transient coronal sigmoid in active region NOAA 11719, which activates and erupts, leading to a fast halo CME. In the following, we summarize the important results of this study: 1. The observations of the active region in the hot EUV channel of AIA (94 Å) reveal the formation of a large S-shaped FR through a sequence of multiple, localized brightenings involving two large pre-existing coronal loop systems. Considering the fact that the bright emission in 94 Å images corresponds to hot plasma ( T ∼ 6 MK), we attribute the multiple
Paper: 2021A&A...649L...6M (fulltext)
panel ). North is on the left-hand side of each plot. The core has been bypassed by both slices. Table 1. Flux densities of the components of 3C 328. Based on the outcome of the two-frequency VLA observations, the inverted-spectrum core must be self-absorbed, and hence very compact. Its VLBA images resulting from the observations mentioned in Sect. 2 are shown in Fig. 2 . At 5 GHz, the core of 3C 328 is point-like, while at 1.5 GHz it has an additional feature elongated southwards. As the same phase reference was used throughout those observations (J1606+3124), the absolute positions of the source at both frequencies were preserved. The core visible at 5 GHz is co-incident with the brightest part of the source in the 1.5 GHz image, and thus this image depicts a core-jet structure where the jet points to the southern lobe. This means that, apart from the lobe asymmetry, there is another similarity that 3C 328 shares with J1211+743, 3C 249.1, and 3C 334: All these AGNs have jets that point towards the lobe that is devoid of a hotspot. This raises the question of whether the light-travel-time-based model elaborated on in Paper II for J1211+743 and then used to interpret 3C 249.1 and 3C 334 in Paper III could also be applied to 3C 328. To find that out, one has to perform a quantitative test using that model since it imposes tight constraints on both the lower and upper limits of the length of the period of quiescence, t q . Fig. 2. VLBA images of 3C 328 at 1.5 GHz ( left-hand panel ) and at 5 GHz ( right-hand panel ). In both images, contours are at 0.1, 0.2, 0.4, 0.8, 1.6, 3.2, 6.4, and 12.8 mJy beam −1 . The projected angular far-side and near-side arm lengths of 3C 328 ( and ) respectively translate to 212 kpc and 232 kpc of the projected linear size, whereas the projected angular size of the jet () translates to 209 pc of the projected linear size. For these linear sizes, the lower () and upper () limits of t q were calculated using Eqs. (2) and (4) of Paper II for β adv = 0.35, three values of β jet (0.8, 0.9, and 1), and five values of the angle between the line of sight and the AGN axis θ . Since 3C 328 is a galaxy and not a QSO, only the values θ ≥ 45° are applicable to it ( Barthel 1989 ). The results are shown in Table 2 . Due to the small size of the jet in 3C 328, the calculated limits are almost independent from the adopted value of β adv . The definitions of β jet and β adv are provided in Paper II. Table 2. Lower and upper limits to the length of the quiescent period. 4. Conclusion Taking into account that for at least two restarted AGNs – 3C 293 ( Machalski et al. 2016 ) and J1835+6204 ( Konar et al. 2012 ) – the values of t q also stay within the limits shown in Table 2 , it is plausible that the above limits for t q may be valid for 3C 328, and as such the light-travel-time-based scenario can be used to explain its nature. It follows that the activity of 3C 328 has restarted and, given the minute size of its jet, that
Paper: 2021A&A...649A.121B (fulltext)
shows the maximum value of the swirling strength, λ max , as a function of z for the three simulations. From it, we notice that in the convection zone the stronger the initial magnetic field, the weaker the swirls. However, the overall behavior is similar and the local peak is reached in all three cases around z ≈ −500 km. Thus, relatively strong vortices can develop in the convection zone, independently of the initial magnetic configuration of the simulation. In the photosphere, a local minimum in swirling strength is reached. It is a result of the slowdown of vortical motions in the rapidly expanding convective overshoots as a consequence of angular momentum conservation ( Nordlund et al. 1997 ). For the two MHD simulations, the minimum is reached in the lower half of the photosphere, while the Hydro simulation reaches its minimum at z ≈ 500 km. This difference can be attributed to strong arches of swirling strength in the hydrodynamical simulation, which permeate the photosphere and extend precisely up to z ≈ 500 km, as it can be seen in Fig. A.1 . These arches are less developed in the MHD simulations, probably as a consequence of the magnetic canopy. Fig. A.1. Volume rendering of the swirling strength seen from the side for three different simulations at arbitrary time instants: ( top ) MHD simulation with an initial vertical, homogeneous magnetic field of 200 G; ( middle ) MHD simulation with an initial field of 50 G; ( bottom ) purely hydrodynamical simulation. The red surface indicates the τ 500 = 1 surface. Swirls with a period larger than 10 min ( λ 2.09 × 10 −2 Hz) have been neglected and values above 1 Hz are saturated. Finally, looking at the photosphere and chromosphere, we observe that λ max increases exponentially for the MHD 50 G simulation, while this growth is only linear for the MHD 200 G simulation. In the Hydro case instead, there is a growth only in the low chromosphere up to approximately z ≈ 1200 km. Above this height, the swirls become weaker again. Clearly, magnetic fields favor the generation of swirls. However, it seems that too strong magnetic fields can also suppress the formation of strong vortices. We surmise that there exists an optimal value of the initial magnetic field for which the strength of the vortices is maximal. Fig. A.2. Distribution of the inclination cos( θ ) of the swirling strength vector as a function of altitude z for three different simulations: ( a ) MHD simulation with an initial vertical, homogeneous magnetic field of 200 G, ( b ) MHD simulation with an initial field of 50 G, and ( c ) purely hydrodynamical simulation. The inclination θ refers to the angle between the swirl axis and the vertical direction. Horizontally oriented swirls are binned close to cos( θ ) = 0, while vertically oriented swirls are close to cos( θ ) = 1. The maximum value of swirling strength λ max at each height z , for each simulation, and with a minimum of 10 occurrences per bin, is shown in panel d . The
Paper: 2020MNRAS.496.3035L (fulltext)
/ L ≈ 0.02. Similar results were found later by Goddard et al. ( 2016 ) with a large array of SDO data. They performed a statistical study of decaying kink oscillations in 120 coronal loops. The best-fitting results were given. In particular, with the mean period P ≈ 350 s and damping time τ ≈ 535 s, the ratio was found τ/ P ≈ 1.53 or γ/ω ≈ 0.10. The average length of analysed loops is L ≈ 240 Mm. We use the above data and the mean density contrast in oscillating coronal loops χ ≈ 0.3, given by Aschwanden et al. ( 2003 ). The aspect ratio, given above, indicates that the fundamental kink oscillations occur at dimensionless axial wavenumber k n a = π a / L ≈ 0.06. Substituting the mean density contrast in equation ( 32 ) we find that Equation ( 50 ) implies that the cut-off wavenumber exceeds the actual observable axial wavenumber of principal kink mode if the source distance d ≤ 70 a , or in terms of loop length if d ≤ 1.4 L . This is a realistic restriction, as the observable distance between the nearby flare sight or erupted filament and the impacted coronal loops is comparable with the loop length. In this case, the externally excited principal kink mode of coronal loops exists in the leaky regime and it is subjected to damping due to the wave leakage effect. With known oscillation parameters ω, γ/ω, χ and loop length L and semiwidth a , we can consider equations ( 43 ) and ( 45 ) as a system for unknown values V Ae and d . We then obtain the following estimates: V Ae ≈ 1250 km s −1 and d ≈ 0.6 L . 5 CONCLUSIONS We have studied the dispersion properties of fast oscillations in a coronal plasma slab with anchored ends. It is assumed that these oscillations are excited by a finitely distant external source. The coronal slab is considered as an enhancement in plasma density with a discontinuous boundary, embedded within uniform environment. A magnetic field is assumed to be everywhere uniform. Using the Fourier transform, the inhomogeneous boundary value problem was solved to obtain the solutions to wave variables in all the regions of the considered model. Corresponding solutions with known coefficients inside the slab contain the so-called dispersion function D (ω) in their denominators and this function completely governs the properties of the slab eigen-oscillations, excited by the external driver. The results of the study are summarized as follows. Both the kink and sausage modes are excited. There are two sets of zeros of the dispersion function. Each of these systems of roots prescribes both axially asymmetric and axially symmetric oscillations. With the appropriate condition that the distance between the source and the slab axis d is much more than the slab semiwidth a (i.e. d ≫ a ), we have shown that the first set of solutions corresponds to kink oscillations of the slab with finite amplitudes and sausage oscillations with negligible amplitudes, proportional to ∝ δ ∝ exp [− 2 m e ( d − a )] ≪ 1, whereas another set prescribes the
Paper: 2021A&A...646A..70S (fulltext)
were likely formed in-situ at the Galactic halo (in black). Interestingly, it seems here that the stars from Gaia Enceladus are the ones closer to the extension of the linear relation, while those that deviate the most are in-situ stars. Fig. 5. Beryllium abundances as a function of metallicity. The points are the same as in the left panel of Fig. 2 , except that we colour the four extremely metal-poor stars found to likely be in-situ halo stars in Fig. 4 in black. In particular, we should remark here the position of stars BPS BS 16968 and CD-33 1173, for which we attempted to derive Be abundances for the first time. These stars are the two open blue triangles plotted over each other at about [Fe/H] ∼ −3.0 and log(Be/H) ∼ −13.8 in Fig. 5 . They are the most important points that demonstrate how stars from Gaia -Enceladus remain closer to the linear relation in a wide range of metallicities. We believe any extended discussion about Be abundances in stellar sub-components of this regime is for now precluded by our small sample. Nevertheless, two main results can be obtained from our analysis: (i) the deviation from the linear relation at [Fe/H] ∼−3.0 is real, although not universal; and (ii) the scatter we observed at these metallicities seems related to the inhomogeneous nature of the early halo. 4.3. Inhomogeneous star formation in the early Galaxy Having established our observational results, we now searched for a possible explanation to the increased scatter of points in the log(Be/H) versus [Fe/H] at the extremely metal-poor regime. To do that, we summarised results from an extensive body of investigations that have looked at the details of the first phases of Galactic chemical enrichment. The interpretation we would like to advance is that the main drivers of the observed scatter are the abundances of Fe and not those of Be. Or, in other words, that at the time of the formation of the early halo, Fe abundances (in a Galactic context) were more inhomogeneous than the Be abundances. 4.3.1. Early enrichment in beryllium The idea that abundances of Be (and of other cosmic-ray spallation products) can have smaller scatter when compared to abundances of elements produced by stellar nucleosynthesis was first suggested by the results of Suzuki et al. (1999) and Suzuki Yoshii (2001) . These authors computed chemical evolution models of an inhomogeneous early Galactic halo. Based on these models, Be abundances were suggested to be a good cosmochronometer for the early stages of the Galaxy ( Beers et al. 2000 ; Suzuki Yoshii 2001 ). This suggestion was tested by Pasquini et al. (2004 , 2007) who derived Be abundances in turn-off stars of the globular clusters NGC 6397 and NGC 6752. These measurements were very challenging, but ages obtained comparing the Be abundances with predictions of chemical evolution models (from Valle et al. 2002 ) were shown to be in excellent agreement with ages derived using isochrones. Suzuki Yoshii (2001) seem to have
Paper: 2016ApJ...823..127N (fulltext)
and quasi-bolometric light curves of the photospheric phase and 56 Ni tail. The sample is heterogeneous, ranging over a factor of 30 in plateau luminosity and 56 Ni mass, and over a range of 2.5 mag/100 days in decline rate. Namely, the sample includes luminous, intermediate, and low-luminosity SNe II as well as SNe that were classified as II-P and as II-L. We analyze this sample, calculating the new observables that we derive here, and find the following results. 1. 56 Ni contribution is significant in most of the sample. In many SNe, it is the source of about 30% of the time-weighted integrated luminosity during the photospheric phase, but there is a single SN (2009ib) in the sample where 56 Ni contributes 72%, and several SNe where it powers about 10%. 2. 56 Ni contribution can possibly flatten and/or extend the photospheric light curve. We find that if flattening were the only effect, then SNe that are observed to have a flat plateau would have shown a decline rate of about 1 mag/100 days in the absence of 56 Ni. If, on the other hand, the only effect of 56 Ni were to extend the plateau, then without 56 Ni, the photospheric phase would have been typically shorter by a factor of about 1.15–1.35 and in the most extreme cases by almost a factor of two. 3. 56 Ni contribution cannot explain the entire range of observed decline rates, which is mostly an intrinsic property of the cooling envelope. Several independent lines of evidence suggest that the dominant effect of 56 Ni is flattening of the light curve. This includes a significant anti-correlation between the level of 56 Ni contribution and the decline rate (i.e., 56 Ni is more important in SNe with a flatter plateau). This is expected if 56 Ni is mixed throughout a significant fraction of the envelope (as suggested by Bersten et al. 2011 for SN 1999em). With this interpretation, the picture that arises from our results is that flat plateaus are not a generic feature of cooling envelope emission. Instead, many observed plateaus include a significant contribution of 56 Ni and without it, many of the SNe II-P would have shown a decline rate of up to 1 mag/100 days. Note that while not being expected a priori, there is no need for a “conspiracy,” so the combination of cooling envelope and 56 Ni emission would result in plateaus. First, SNe II exhibit a range of decline rates, thus in some cases the combined emission produces a plateau and in others it does not. Second, it seems that even without 56 Ni, the cooling envelope emission often shows a moderate decline rate where the luminosity drops by less than a factor of two over 100 days. In these cases the contribution of 56 Ni is enough to make the light curve even flatter. The coincidence in SNe II seems to be the fact that the contribution of the cooling envelope and 56 Ni-powered emission is comparable over a wide range of light curve luminosities. This is a result of the well known tight correlation between the luminosity and 56 Ni
Paper: 2016ApJ...827L..27P (fulltext)
because the gas is ionized. We searched for a warm ionized gas associated with GS040 using the Wisconsin H α Mapper Northern Sky Survey (WHAM-NSS) data (Haffner et al. 2003 ). The survey has an angular resolution of 1°, and, around the GS040 area, it provides H α spectra covering from −85 to +100 km s −1 at spectral resolutions of 12 km s −1 . We have examined the H α intensity map integrated over the velocity range of GS040 (−85 to −66 km s −1 ), but could not detect an associated emission ( 0.05 R where ). Note that this area is bright in the total H α intensity map with a mean intensity of about 3 R , so that we do not expect to see the faint emission associated with either GS040 or CHVC040 in the all-sky H α maps (Dennison et al. 1998 ; Finkbeiner 2003 ). We also searched for an associated hot ionized gas using the 0.1–2.4 keV image of the ROSAT All-Sky X-ray Survey (1 pixel scale = 44′′; Voges et al. 1999 ) but could not detect any emission. CHVC040 belongs to the “Galactic Center Negative” (GCN) HVC complex, which is a collection of small discrete HVCs sparsely distributed over a area within and to 10° (Wakker van Woerden 1991 ; Winkel et al. 2011 ). A kinematic model has been proposed where GCN is a smooth gas flow starting at at a heliocentric distance of 35 kpc and crossing the Galactic plane obliquely at 15 kpc (Jin 2010 ). For comparison, GS040 is probably located near the Scutum-Centaurus (Sct-Cen) arm at a distance of ∼20 kpc (Dame Thaddeus 2011 ) because the disrupted interstellar medium (ISM) is seen only at velocities below that of the Sct-Cen arm ( km s −1 ), not at higher velocities (see Figures 2 and 3 ). We examined lists of known stellar objects, H ii regions, OB stars, and SN remnants, but no known sources are likely to be associated with GS040. There is one H ii region (G039.864+00.645) in the WISE catalog of Galactic H ii regions (Anderson et al. 2014 ) that is located at west of the hub. But this H ii region has a systematic velocity of −40.9 km s −1 (Anderson et al. 2011 ) and is enclosed by a small (68′′) dust shell, so it cannot be responsible for the H i shell. Instead, the distance of ∼20 kpc is not unreasonable for CHVC040, because GCN does not have a smooth extended envelope like other HVC complexes, and it appears to be composed of several subpopulations that do not share a common origin (Winkel et al. 2011 ). Note that the geometrical center of GS040 is well below the Galactic plane ( pc where ), i.e., at while the midplane there is at because the Galactic plane is warped in the outer Galaxy (Levine et al. 2006 ). It is difficult to imagine the SN origin for a supershell at such height, and CHVC040 is most likely the energy/momentum source for GS040. The total energy deposited ( E E ) in the Galactic disk by CHVC040 can be inferred from the parameters of the GS040 supershell. The radius of GS040 is pc, while its mass at km s −1 is including the cosmic abundance of helium. If we account for the mass
Paper: 2021MNRAS.502.2251G (fulltext)
the posteriors estimates of the TNM model is shown by the red line. Right-hand panel: V φ distribution for mono-age stellar particles. Colour code as in Table 2 . Table 4. TNM model parameters: w i is the mixture weight of the i -th component (i.e. young disc, old disc/thick disc, halo), while 〈 V φ 〉 and are the corresponding mean and the standard deviation of the normal distribution. Note that 0 ≤ w i ≤ 1 and ∑ i w i = 1. We also tested alternative models with a different number of components. The results for a double Gaussian distribution such as halo+disc or disc+disc are not statistically significant and do not provide good curve fitting of the data. A four-component model, aimed to represent additional disc populations, produced non-significant improvements to the model. We note that the weights of the two disc components listed in Table 4 correspond to the total mass of the young disc ( Gyr) and of the old disc/thick disc (), respectively (cf. Table 2 considering the different vertical cut applied). The comparison with the mono-age distributions shown in Fig. 7 (right-hand panel) confirms that these two disc components represent the superposition of the stellar generations formed in the age intervals 0–4 and 4–10 Gyr. Note that the old thin disc and the thick disc populations are difficult to distinguish because of their very similar velocity distributions. Moreover, the difference between the mean rotation velocity of the AqC4 discs is This value is smaller than in the MW, where (Han et al. 2020 ). Actually, this difference depends on the fact that in AqC4 the two main disc components represent the young and the old/thick discs instead of the whole thin disc and the thick disc, as in the MW. Moreover, the high rotation velocity of the AqC4 disc is consistent with its more compact structure with respect to the MW disc (Section 3.1 ). Finally, the halo component shows a small prograde rotation as observed in the inner halo of the MW, while its velocity dispersion km s −1 is about 70–100 per cent higher than in the solar neighbourhood (Re Fiorentin et al. 2015 ; Bland-Hawthorn Gerhard 2016 , and references therein). In summary, here we robustly support what is discussed in Section 3 : in the SSR, we can distinguish at least two kinematic disc components which rotate faster than what is measured in the solar neighbourhood (as discussed globally in Section 4.1 ), and with a smaller discrepancy between the mean values. These results extend the spatial characterization of AqC4 stellar disc beyond the determination of the scale lengths and scale heights, allowing us to properly investigate the complete phase-space of mono-age populations, as usually done in Galactic surveys. 4.3 Kinematics of mono-age disc populations Fig. 8 shows the three median velocity components , , and the dispersions (, , ) as a function of R , computed in radial annular bins of Δ R = 0.25 kpc, for the mono-age stellar disc populations younger than 10 Gyr and with | Z |