adsabs/FOCAL huggingface dataset:
To investigate the particle precipitation or transport sites during the flare, we used HXR 25–50 keV and NoRH 17/34 GHz images. We chose the Pixon algorithm (Metcalf et al. 1996) for the RHESSI image reconstruction. The Pixon method is considered to be the most accurate algorithm (Hurford et al. 2002). The integration time for each image was 20 s. We utilized NoRH 5 s cadence intensity images (R+L) at 17 and 34 GHz. Figure 4 displays the HXR 25–50 keV (blue) and NoRH 17 GHz (red) contours overlaid on the AIA 1600 Å images at ∼23:51 UT. These images are used at the peak time of each of the bursts observed in the HXR 25–50 keV and microwave channels. Figure 4(a) shows a small filament rising at the flare site. We can easily identify the two legs of the filament (marked by N and S in the figure). Generally, the filaments are observed in the chromospheric images (e.g., Hα and AIA 304 Å). However, if the kink-unstable filament is heated during magnetic reconnection, it is often observed in the AIA 1600 Å channels (e.g., Kumar et al. 2012 and Kumar & Cho 2014). The locations of both the 25–50 keV sources and the 17 GHz source are almost at the quasi-circular ribbon. However, their centroids constructed at the 90% level of the peak intensity are not cospatial. We note that the 25–50 keV sources (centroid) are located close to the legs of the filament, whereas the 17 GHz source is located at the northern leg of the filament. It seems that these sources are at the footpoints of an underlying flare loop. To identify the location of the flare loop, we selected the AIA 94 Å hot channel image at 23:53 UT. In this image (panel b), we show the NoRH 34 GHz contours (yellow) overlaid on AIA 94 Å images at ∼23:53 UT. This is done to show the coronal loops associated with the eruption of the small filament. The overlying 34 GHz source is cospatial with the small loop located above the quasi-circular ribbon. This 34 GHz emission may be the evidence of trapped nonthermal electrons in the loop. The rising filament is heated during reconnection with the ambient fields and then decays into an untwisting jet, as seen in Figures 4(b)–(e). Furthermore, during ∼23:53–23:59 UT, we see the footpoint HXR sources. The NoRH 17 GHz sources cover the quasi-circular ribbon (Figures 4(c)–(d)) and could be the emission from the trapped electrons in the loop. This suggests that the bursts are most likely caused by the same population of nonthermal electrons with different energies.
Citation: Kumar et al. 2012
Citation function: Uses
The contribution to the submm background from LBGs is still poorly constrained; however, our data can finally address this question, since we now have a robust detection of the average 850 μm flux density of LBGs at three epochs, at least those with ultraviolet luminosities of L1700 ≈ 1029 erg s−1 Hz−1, characteristic of galaxies in our sample (Table 2). Using our average stacked 850 μm flux densities, we estimate surface brightness densities of 1700, 600, and 100 mJy deg−2 of LBGs at z ∼ 3, 4, and 5, respectively. By comparison, the total background at 850 μm inferred from COBE-FIRAS is 3.1–4.4 × 104 mJy deg−2 (Puget et al. 1996; Fixsen et al. 1998; Hauser & Dwek 2001; Lagache, Puget & Dole 2005). Summing these separate surface brightness densities together, we find that the LBGs with L1700 > 1029 erg s−1 Hz−1 in our z ∼ 3, 4, and 5 samples comprise around 6–8 per cent of the submm background at 850 μm (where the range of values simply reflects the uncertainty in the COBE-FIRAS result). However, the true contribution from LBG-like galaxies to the submm background will come from a wider range in redshift, not just from the rather narrow redshift slices we have sampled (see Fig. 1), and from sources falling out of our samples due to incompleteness. To determine the total (corrected) contribution from LBGs over 3  z  5, we assume that LBGs have a constant comoving number density (which is a reasonable assumption since the bright end of the luminosity function for LBGs shows little evolution over 3  z  5; Bouwens et al. 2007; McLure et al. 2009; Reddy & Steidel 2009). Starting with the z ∼ 3 sample (which is our most complete subsample of LBGs), we can integrate over the comoving volume element for this redshift range, scaling the LBG submm background contribution accordingly. We find that the total contribution to the submm background from LBGs over the redshift range 3  z  5 is likely to be closer to 14–20 per cent. This result is consistent with Webb et al. (2003), who estimated an upper limit to the contribution to the submm background from 1  z  5 of less than 20 per cent.
Citation: Hauser & Dwek 2001
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The UK Chemistry and Aerosol (UKCA) model (Morgenstern et al. 2009; O’Connor et al. 2014; Archibald et al. 2020) is a framework that we use to describe the global atmospheric chemical composition of our simulated exoplanet. UKCA includes aerosol and gas-phase chemistry and is coupled to the UM dynamics. It uses the UM components for large-scale advection, convective transport, and boundary layer mixing of its aerosol and chemical tracers (O’Connor et al. 2014; Archibald et al. 2020). UKCA contains a large number of gas-phase and heterogeneous chemical reactions, some of which we have included in our chemical network. Furthermore, the chemistry schemes in UKCA describe wet and dry deposition (Giannakopoulos et al. 1999). In this study, we use the Stratospheric (Strat, Morgenstern et al. 2009) and Stratospheric-Tropospheric (StratTrop, Archibald et al. 2020) chemistry schemes. Originally, StratTrop includes 75 chemical species that are connected by 283 reactions (Archibald et al. 2020). We used a reduced version of the UKCA chemistry schemes (Table 2) to quantify the impact of the different chemical mechanisms on the atmospheric chemistry of a tidally-locked exoplanet. First, we use a simple network that describes the Chapman mechanism of ozone formation (Chapman 1930), following Yates et al. (2020). Second, we add the reactive hydrogen (HOx) catalytic cycle, where HOx denotes the ensemble of atomic hydrogen (H), the hydroxyl radical (OH) and the hydroperoxy radical (HO2). We include this cycle to account for ozone chemistry following the oxidation and photolysis of water vapour. Lastly, we add the nitrogen oxide (NOx) catalytic cycle, including NO and nitrogen dioxide (NO2) to the network. We also include other oxidized nitrogen species, such as nitrate (NO3), nitrous oxide (N2O), and the reservoirs nitric acid (HNO3), and dinitrogen pentoxide (N2O5). Collectively, these nitrogen species belong to the NOy family and can also influence ozone chemistry. In our simulations, lightning is the main source of NO that initiates further NOy chemistry, as described in Section 2.4. In the upper atmosphere, the slow termolecular reaction between N2 and O(1D) provides another source of NOy, but this does not impact the lightning-induced chemistry that occurs at altitudes below 20 km.
Citation: Strat, Morgenstern et al. 2009
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From the observational side, the deep FDS data further confirm that the bulk of the gravitational interactions between galaxies happened on the W-NW sub-clump of the cluster. In fact, this is the only region of the cluster, inside the virial radius, where the intra-cluster baryons (diffuse light and GCs) are found, that is, the bridge between NGC 1399 and FCC 184 (Iodice et al. 2016), the intra-cluster light and GCs between FCC 184, FCC 170, and FCC 161 (Iodice et al. 2017b), and the new faint filaments between FCC 143 and FCC 147 (see Sect. 5.1 and Fig. 13). The gravitational interactions could have also modified the structure of the galaxy outskirts and produced the intra-cluster baryons. Compared with simulations by Rudick et al. (2009), the diffuse form observed for the ICL is consistent with the scenario where this component formed by stripped material from the outskirts of a galaxy in a close passage with the cD (Iodice et al. 2017b). In this area of the cluster, the stellar envelope of some ETGs is asymmetric, appearing more elongated and twisted in one direction, while the outskirts of galaxies at larger distances from the cluster centre have a more regular shape. Mastropietro et al. (2005) show how harassment can induce twists in the outer isophotes of dwarf galaxies. More massive and luminous galaxies, like the elliptical galaxies in this region of the Fornax cluster, have a deeper potential well, therefore the stripping of stars by harassment implies even stronger tidal forces. A disturbed morphology in galaxy outskirts could also result from the ongoing accretion of smaller satellites. Simulations on the mass accretion and stellar halo formation for different stellar masses (1010 − 1013 M⊙) show that the outskirts of galaxies (for μr ∼ 27 − 31 mag arcsec−2) appear with a quite disturbed morphology and with an overall elongated shape (Michel-Dansac et al. 2010; Cooper et al. 2015; Monachesi et al. 2018). The structure of the stellar envelope, as well as the shape of the SB profile, depends on the mass and number of the accreted progenitors. The ETGs showing asymmetric and diffuse envelopes (FCC 161, FCC 167, FCC 184, see Sect. 5.1 and Fig. 13) are in range of stellar mass (∼3 − 10 × 1010 M⊙) comparable with simulations and they could still build up their envelope. As noticed by Iodice et al. (2017b), a fraction of the ICL population in this region of the cluster could also come from lower mass dwarf galaxies that are tidally disrupted in the potential well of the massive galaxies, which are therefore contributing to the mass assembly in their halo. This is further supported by a recent study from Venhola et al. (2017), based on FDS data, that found a drop in the number density of LSB galaxies at cluster-centric distances smaller than ∼180 kpc.
Citation: Rudick et al. (2009)
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At intermediate energies, Fig. 1 shows that non-thermal bremsstrahlung processes become more important. We find this dominates the keV X-ray emission during the first ∼2 Myr. While OY22 demonstrated this non-thermal emission would likely not be detectable in external galaxies, X-ray observations towards the Galactic Fermi bubbles reveal substantially higher emission than that computed with our model at a similar bubble age (see the comparison in OY22; also Snowden et al. 1997; Kataoka et al. 2013; Predehl et al. 2020). Although this may suggest better X-ray detection prospects for external galaxy bubbles than our results would imply, we note that these Galactic observations may include a very significant thermal bremsstrahlung contribution from all the gas in the Milky Way halo, which likely extends to a radius of ∼250 kpc (i.e. far larger than the size of our simulation box, see e.g. Blitz & Robishaw 2000; Grcevich & Putman 2009), the Galactic bulge, and gas heated by the shocks associated with the bubble (Yang et al. 2022; see also Zhang & Guo 2021), none of which is included in our model (and could evolve as the bubble ages). Moreover, features external to the Galactic bubbles (e.g. the North Polar Spur; Kataoka et al. 2013) are not included in our model, but may make a contribution to the observed X-ray emission. We thus consider a direct comparison between the X-ray emission from our model and that observed from the Galactic Fermi bubbles/halo to be complicated by substantial thermal emission and emission from structures not associated with the bubbles. Addressing these additional contributions is non-trivial, and is not necessarily informative to predict the observational prospects of bubbles around external galaxies where thermal emission contributions could differ greatly. After the first few Myr, non-thermal bremsstrahlung X-rays are surpassed by inverse Compton emission, which dominates the X-ray emission from the bubble by 7 Myr. At these later times, a significant non-thermal bremsstrahlung component can also be seen to emerge in TeV γ-rays. This is attributed to the concentration of gas and CR energy density near the top of the bubble (cf. Section 3.3), although we note that the physical strength of this emission may not be properly resolved.20 This will be explored in more detail in future work (see also Section 4.3).
Citation: Yang et al. 2022
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In this paper, we are focusing on the association of the acceleration and release of SEPs with different types of CME–CME interaction. In fact, the particle acceleration process is a complex one involving many factors, e.g., (1) the CME (or shock) speeds (Richardson et al. 2015; Papaioannou et al. 2016; Kouloumvakos et al. 2019), (2) the shock parameters, e.g., Mach numbers, compression ratios, and shock geometry (Lee 1983; Kozarev et al. 2015; Kouloumvakos et al. 2019), (3) the magnetic connectivity between the spacecraft and the source (Richardson et al. 2014; Xie et al. 2019), (4) the level of seed particles, and (5) the level of turbulence (e.g., Wanner & Wibberenz 1993; Laitinen et al. 2018; Strauss & le Roux 2019). Figure 2 shows a relationship between the CME 3D speed and the SEP Ip in NP3D group, which is consistent with the results in Kouloumvakos et al. (2019). The penultimate point has been widely discussed (e.g., Gopalswamy et al. 2004; Kahler & Vourlidas 2014; Ding et al. 2015), involving the role of the flare material and the interacting CMEs on seed particles. In Ding et al. (2015), the authors used a Fe/O ratio of 2.0 as the threshold for the presence of flare material and found that all events except one with Fe/O > 2.0 are in the “twin-CME” scenario, indicating the presence of flare seed material that are possibly from pre-flares. In Kahler & Vourlidas (2014), they proposed that higher SEP Ip could be explained by increases in both CME rates and seed particles during times of high solar activity instead of being due to CME–CME interaction. In Figure 7, the higher Ip in the cases of (nearly) interacting CMEs gives a hint that more seed particles will be accelerated if the priCME eruption is close to the preCME eruption, but it is still necessary to figure out which factor, i.e., the solar activity, the flare, or the preCME, mainly controls the level of the seed populations. The launch of the Parker Solar Probe (PSP) can contribute to this investigation in the future, because the spacecraft will reach as close as 8.86 Rs from the solar surface by 2024. During its first two perihelion passages, a series of SEP events have been studied, focusing on the acceleration mechanisms and seed population preconditioning (see McComas et al. 2019; Schwadron et al. 2020).
Citation: Gopalswamy et al. 2004
Citation function: Background; Background
We consider that the amplified hotspot magnetic field B is turbulent, and that the large-scale background field downstream of the reverse shock is Bjd, nearly perpendicular to the shock normal because the perpendicular component is compressed and enhanced by a factor of 4 to 7 (i.e. Bjd ∼ rBj). In such a case, to accelerate particles up to an energy Ec via a diffusive mechanism, the mean-free path λc ∼ rg(γc, B)2/s in the shock downstream region, where B is a small-scale field, has to be smaller than Larmor radius in Bjd (Lemoine & Pelletier 2010; Reville & Bell 2014).4 The condition λc ≲ rg(γc, Bjd) is satisfied when the magnetic-turbulence scalelength is

(23)

\begin{equation}
s \ge \frac{E_{\rm c}}{eB}\left(\frac{B_{\rm jd}}{B}\right) = r_{\rm g}(\gamma _{\rm s},B),
\end{equation}

where rg(γs, B) is the Larmor radius of protons with energy

(24)

\begin{eqnarray}
E_{\rm s} &amp;=&amp; E_{\rm c}\left(\frac{B_{\rm jd}}{B}\right) = 0.07 E_{\rm c}\left(\frac{r}{7}\right) \left(\frac{B_{\rm j}}{\rm \mu G}\right) \left(\frac{B}{100 \,\rm \mu G}\right)^{-1}\nonumber\\
&amp;&amp;\sim 10 \left(\frac{r}{7}\right) \left(\frac{\nu _{\rm c}}{10^{14}\,{\rm Hz}}\right)^{\frac{1}{2}} \left(\frac{B_{\rm jd}}{\rm \mu G}\right) \left(\frac{B}{\rm 100\,\mu G}\right)^{{-}\frac{5}{2}}\,\,{\rm GeV},
\end{eqnarray}

where we take B ∼ 100 μG and Bj ∼ μG as characteristic values. Note that

(25)

\begin{equation}
\frac{s}{\rm cm} > 5\times 10^{11}\left(\frac{r}{7}\right) \left(\frac{\nu _{\rm c}}{10^{14}\,{\rm Hz}}\right)^{\frac{1}{2}} \left(\frac{B_{\rm jd}}{\rm \mu G}\right) \left(\frac{B}{\rm 100\,\mu G}\right)^{{-}\frac{5}{2}}
\end{equation}

is greater than c/ωpi in equation (14), as required. Note however that this limit, s ≳ 500 c/ωpi for typical values considered in this paper, cannot be fulfilled by Weibel-generated turbulence with scale ∼c/ωpi. Therefore, the maximum energy achieved by electrons in the jet reverse shock, Ec, cannot be constrained by Weibel instabilities.
Citation: Reville & Bell 2014
Citation function: Uses
Note that the evolution scenario in Fig. 2 is qualitatively different from previous expectations that assumed standard viscous damping and a large saturation amplitude, see e.g. Levin (1999). There it was proposed that sources slowly spin-up at low temperatures outside of the instability region (where cooling becomes slow) followed by quick r-mode heating once the source enters the instability region and similarly fast spin-down and cooling segments which complete a cycle. There would then be no gravitational wave emission from known radio pulsars since sources would leave the instability region very quickly (Levin 1999). However, this scenario is not compatible with the well-established hypothesis that LMXBs are the progenitors of millisecond pulsars, because their observed large temperatures and frequencies (Haskell et al. 2012) place them firmly inside the instability region, as shown e.g. in fig. 1 of Alford & Schwenzer (2013). These astrophysical observations are furthermore theoretically explained by deep crustal heating due to pycnonuclear reactions (Brown, Bildsten & Rutledge 1998). For large r-mode saturation amplitudes LMXBs either could not spin-up to the high frequencies of observed radio pulsars or would be spun down very quickly, which should rule out this scenario. A modified scenario that might account for the observed spin limit of pulsars would be the presence of enhanced damping that does not increase with temperature, for instance due to a viscous boundary layer (Andersson et al. 2000) or mutual friction in a superfluid and superconducting core (Haskell, Andersson & Passamonti 2009; Haskell et al. 2012). In this case r-modes could be completely damped up to the frequencies of observed radio pulsars and it would be very unlikely that any of the known pulsars currently emits gravitational waves (Andersson et al. 2000). However, estimates on mutual friction are still rather uncertain (Haskell et al. 2009, 2012), and we have recently shown that even in the benevolent scenario of a thin viscous boundary layer the damping is very likely not sufficient to explain the fastest spinning sources (Alford & Schwenzer 2013). In Fig. 2 we show the case that crustal heating dominates r-mode heating whereby a source spins up along a vertical trajectory. If r-mode heating dominates, a source would spin-up along roughly the same trajectory along which it subsequently spins down. Another option is that the spin-up stalls since a steady state is reached where the accretion spin-up is balanced by r-mode spin-down (Wagoner 2002). However these different scenarios during the LMXB phase lead to the same qualitative evolution once the accretion ends.
Citation: Brown, Bildsten & Rutledge 1998
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Finally in this paper, we restricted ourselves to a very limiting case when the ISM is at rest with respect to the star. This two-jet solution can, in principle, be generalized by adding the interstellar flow. Let us consider an arbitrary plane perpendicular to z-axis. This plane cuts a circle from the astropause. In the case of subsonic ISM flow, we can consider planar potential solutions around circles for each plane. According to the d'Alembert paradox, the force acting on each circle is zero. Therefore, the tube of the astropause should not be deflected into the tail, although the circle could be deformed to the ellipsoidal shape in the self-consistent solution. This scenario works, if we consider the interstellar flow to be ideal and incompressible. However, numerical results (e.g. Opher et al. 2015) show some bending of the jets towards the tail. This bending in numerical models (for slow incompressible ISM flows) is connected with the numerical dissipation inherent in the numerical schemes. Numerical viscosity may cause the boundary layer breakage on the surface of the astropause, which consequently causes the pressure asymmetry that deflects the astropause. This may be the explanation for the fact that the tube of the astropause is always deflected to the tail in the numerical models. The above-described numerical effects have nothing to do with physical dissipation phenomena responsible for the bending of real astrospheres. The physical dissipation mechanisms (e.g. magnetic reconnection, finite resistivity, Hall effects) should be explored as a possible cause of the astropause bending in the models with slow subsonic ISM flow. For the fast supersonic ISM flow, the BS is formed around the astropause. The post-shock ISM flow is vortical, and the d'Alembert paradox does not work in this case. Therefore, the bending of the astrospheric jets into the tail direction is easier to justify for the supersonic relative ISM/SW motion.
Citation: Opher et al. 2015
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Using the whole sample (\hbox{$\bar{z}=0.40$}z̅ = 0.40), we find a decreasing faint end for both datasets with consistent values between HST (α =  − 0.76 ± 0.07) and Subaru (α =  − 0.78 ± 0.06). Separating between low-redshift (\hbox{$\bar{z}=0.29$}z̅ = 0.29) and high-redshift (\hbox{$\bar{z}=0.51$}z̅ = 0.51) samples, we find an evolution of the faint end slope of 1.7σ with HST and 2.6σ with Subaru. There is thus a mild decrease of the faint end slope (less negative α) with increasing redshift over the range (0.187 z 0.686). This evolution is in good agreement with recent papers in the literature: in particular Zenteno et al. (2016) found a decrease of the RS faint end at 2.1σ for a wider range of redshifts (0.1 z 1.13), but with ~ 80% of their clusters being in the same redshift range as ours. De Propris et al. (2013) claim that the evolution in the faint end slope has a significant contribution from surface brightness selection effects. They used HST data of differing depths on a single cluster (MS 1358.4+6254) to show that surface brightness selection effects become important above the formal magnitude limit of their data and that they affect the RS GLF at magnitudes z ≥ 24.5 for 2.7 ks HST exposures (see their Fig. 18). The faint RS for their cluster has F814W − z = 0.25, implying that the SB selection effects in their sample become important at F814W> 24.75. On the other hand, our CLASH data are significantly deeper than theirs (4.1 ks) and we limit our GLFs at F814W 24.5. Therefore, the real SB selection effects noticed in De Propris et al. (2013) should not be playing a role in our space-based results. In addition, De Propris et al. (2013) claim that previous estimates of the evolution in the RS GLF (e.g.,  De Lucia et al. 2007; Rudnick et al. 2009) were also due to SB effects. Both of those works were based on the same ground-based data with a formal magnitude limit of I = 24 or 24.5 (for the low- and high-redshift clusters, respectively) and the evolution in the GLF was seen over the faintest 2 mag. We cannot directly address the role of SB effects in the EDisCS results without detailed simulations on those data (see below for such simulations for our clusters) but the similarity between our HST and Subaru GLFs imply that the EDisCS evolution in the GLF is not dominated by SB effects.
Citation: De Propris et al. (2013)
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