adsabs/FOCAL huggingface dataset:
The coronal mass ejections (CMEs) are frequent discharge of huge energy and massive magnetized plasma from the solar corona into the heliosphere. They are of paramount importance in space physics for their key role in extreme space weather and geo-effectiveness, e.g. (Gosling 1993; Low 2001; Schrijver & Siscoe 2010; Cannon et al. 2013). In last few decades, the understanding of CMEs improved significantly because of space and ground-based observational data with the help of various modelling efforts. The studies are focused on the morphological and kinematic evolution of CMEs in the heliosphere, e.g. (Lindsay et al. 1999; St Cyr et al. 2000; Zurbuchen & Richardson 2006; Chen 2011; Webb & Howard 2012; Wang et al. 2016; Lugaz et al. 2017). By considering the number of CMEs emitted from the Sun during solar maximum and variations in their respective speeds, the interaction between multiple CMEs in the heliosphere is expected to be more frequent. The collision of multiple CMEs highly affect their dynamic evolution properties and contribute to enhanced geo-effectiveness, e.g. (Wang, Wang & Ye 2002; Farrugia & Berdichevsky 2004; Lugaz, Manchester IV & Gombosi 2005; Wang et al. 2005; Xiong et al. 2007; Shen et al. 2011, 2012; Lugaz et al. 2012; Temmer et al. 2012). To predict space weather effects near the Earth, an accurate estimation of arrival time of CMEs at the Earth is crucial (Mishra & Srivastava 2014). Besides this, the study of CME–CME and CME–solar wind interactions provide unique observational evidences to understand energy dissipation of large-scale magnetic clouds in interstellar medium and authenticate the physical processes predicted theoretically. Therefore, interaction of multiple CMEs needs to be examined in detail. The various results obtained from studies have justified CME–CME collision as an inelastic/elastic collision or superelastic collision, e.g. (Lugaz et al. 2012; Shen et al. 2012, 2016; Lugaz et al. 2017; Mishra et al. 2017). The magnetohydrodynamics (MHD) numerical simulations have striven to understand the physical mechanism involved in CME–CME interaction, CME–CME driven shock interactions and their consequences, e.g. (Niembro et al. 2015; Jin et al. 2016; Shen et al. 2016; Wu et al. 2016).
Citation: St Cyr et al. 2000
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An example of these implications is the so-called ‘dust budget crisis’ introduced in Section 4.4: the dust masses currently estimated at z > 5 are not compatible with standard dust production channels and require an overhaul in our models of the initial mass function for star formation, of supernova production rates, or of dust growth in the ISM. Overall, the dust production rate would need to increase by one to two orders of magnitudes, as shown by Rowlands et al. (2014). The growth of dust grains through accretion in the ISM has been proposed as a solution (e.g. Mancini et al. 2015; Michałowski 2015; Popping et al. 2017), but there are doubts on the efficiency of accretion at high z, where high dust temperatures due to the CMB (see Section 3.3) keep the desorption time-scale for accreted materials short (Ferrara et al. 2016). The dust budget crisis is not only a problem at high redshift; it is observed, e.g. in the Magellanic Clouds (SMC, LMC). As explained in Srinivasan et al. (2016) using the dust mass fits by Gordon et al. (2014), the dust replenishment time-scale in the SMC from stellar sources alone is expected to be larger than the dust destruction time-scale and, in the worst-case scenario, longer than the lifetime of the Universe. Similarly, the ratio between the best LMC dust mass estimate by Gordon et al. (2014) and the dust injection estimates by Riebel et al. (2012) results in an LMC replenishment time-scale of 34 ± 8 Gyr, exceeding the age of the Universe. Both the high redshift and the local Universe, therefore, show a dust budget crisis that could be alleviated – and, in the best case scenario, fully resolved – if the actual dust masses turned out to be lower than currently estimated, as our results suggest. More specifically, Rowlands et al. (2014) mention that dust opacity needs to be increased by just a factor of 7 to solve the high-redshift crisis (provided dust destruction by SNe is not efficient); in the LMC, the aforementioned replenishment time-scale would decrease to less than 2 Gyr if the dust mass were decreased by a factor of 20.
Citation: Popping et al. 2017
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Citation: Michałowski 2015
Citation function: Motivation
Compared to various studies based on in situ measurements of MFR structures after eruption, the origination of CME-MFRs before and during eruptions remains elusive due to the complex environment in the solar source region and limited observations. At the present time, there are certain hypotheses on the formation process of MFRs. Some studies indicate that MFRs could exist prior to the eruption. For example, both Fan (2001) and Magara (2004) reported findings from magnetohydrodynamic (MHD) simulation that a twisted MFR initially formed below the photosphere can partially emerge into the low corona by magnetic buoyancy. Other studies suggest that the presence of preeruptive MFRs is not necessary and MFRs could be built up in the corona via magnetic reconnection processes associated with flares (Amari et al. 2003; Moore et al. 2001; Antiochos et al.1999; Jiang et al. 2021a, 2021b). To understand the physical processes more precisely for flare–CME events, extensions of the standard 2D flare model have been proposed to account for much broader ranges of quantitative measurements with three-dimensional (3D) features intrinsic to realistic solar eruptions (Longcope et al. 2007; Aulanier et al. 2012; Priest & Longcope 2017; Aulanier & Dudík 2019). For example, quasi-3D models have been developed with a nonvanishing magnetic field component along the axis of the MFR and to illustrate the scenario that sequential reconnection along the magnetic polarity inversion line (PIL) forms the MFR in the first place (van Ballegooijen & Martens 1989; Longcope et al. 2007; Schmieder et al. 2015). This scenario has been widely applied to infer and interpret magnetic reconnection properties based on the observed flare ribbon morphology (Qiu et al. 2002, 2004, 2010; Hu et al. 2014; Kazachenko et al. 2017; Zhu et al. 2020). From such analyses, Qiu et al. (2004) illustrated that there is a temporal correlation between the magnetic reconnection rate and the acceleration of the CME (considered as the eruptive MFR) in the low corona. Such a correlation has been further established by Zhu et al. (2020) based on a statistical study of ∼60 events. In addition, Qiu et al. (2007) and Hu et al. (2014) showed a correlation between the magnetic reconnection flux and the flux contents of the corresponding ICME/MC flux ropes based on modeling results employing in situ spacecraft measurements. These results support the hypothesis that CME-MFRs can be formed by magnetic reconnection during the corresponding flare process. Recent simulation results also indicate clearly that the reconnection flux contributes to the axial (toroidal) flux of the CME-MFR in the early stage (Jiang et al. 2021a; Inoue et al. 2018).
Citation: Moore et al. 2001
Citation function: Motivation
Although it is still unknown why the disc fragmentation for the primordial cases is well described by the simple relation such as equation (1), an important fact is that a barotropic EOS with γeff ≃ 1.1 approximately represents the gas thermal evolution during the cloud collapse for $n \lesssim 10^{19}\, \mathrm{cm}^{-3}$ (Omukai & Nishi 1998). In fact, several simulations study the disc fragmentation assuming the same barotropic EOS with γeff = 1.1 for n ≤ nth, resulting in the evolution described by equation (1) (Susa 2019). Since nth is only the characteristic quantity for this case, the simple scaling of equation (1) may be convincing. This suggests that the disc fragmentation with a different EOS should provide different evolution of Nc,b. For instance, it is well known that adding a tiny amount of heavy elements and dust grains alters the EOS of a collapsing cloud (e.g. Omukai 2000; Bromm et al. 2001; Omukai et al. 2005; Schneider et al. 2006, 2012a; Smith, Sigurdsson & Abel 2008; Jappsen et al. 2009; Safranek-Shrader, Milosavljević & Bromm 2014; Chiaki et al. 2015; Chiaki, Yoshida & Hirano 2016). While previous studies demonstrate that the dust cooling enhances the fragmentation during the cloud collapse (e.g. Meece, Smith & O’Shea 2014; Smith et al. 2015; Chiaki & Wise 2019), its effect on the disc fragmentation remains to be further explored. Tanaka & Omukai (2014) investigate the evolution of the circumstellar disc in metal-poor environments developing 1D semi-analytical models. They predict that the discs with $Z \sim 10^{-5}\!-\! 10^{-3}\, \mathrm{Z}_{\odot }$ are subject to the efficient dust cooling and are more unstable than those for the primordial cases. Machida & Nakamura (2015) consider the disc fragmentation with various metallicities $0 \le Z \le 1\, \mathrm{Z}_{\odot }$, performing a suite of 3D numerical simulations. They find qualitative differences between the cases with $Z \lesssim 10^{-4}\, \mathrm{Z}_{\odot }$ and $Z \gtrsim 10^{-4}\, \mathrm{Z}_{\odot }$; the vigorous disc fragmentation only occurs for the former metal-poor cases. Whereas Machida & Nakamura (2015) use the metallicity-dependent barotropic EOS, Chiaki & Yoshida (2020) recently report 3D simulations of the disc fragmentation solving the energy equation with relevant thermal processes coupled with a non-equilibrium chemical network. They find that for the metal-poor cases with $Z \le 10^{-3}\, \mathrm{Z}_{\odot }$, the disc fragmentation does not necessarily prevent the mass growth of the most massive protostar as many clumps are short-lived owing to the frequent merger or tidal disruption events. Although the above studies suggest the metallicity-dependence of the disc fragmentation, they both only follow the short-term evolution for $\sim 100\, \mathrm{yr}$ since the first emergence of a protostar.
Citation: Chiaki, Yoshida & Hirano 2016
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Comparison with z ≈ 2−4galaxies. High-ionization UV lines have been detected in the spectra of z ≳ 2 galaxies through gravitational lensing (e.g., Stark et al. 2014; Patrício et al. 2016; Vanzella et al. 2016, 2017; Berg et al. 2018), spectral stacking (e.g., Nakajima et al. 2018b; Rigby et al. 2018; Saxena et al. 2020), and deep spectroscopic observations (e.g., Erb et al. 2010; Maseda et al. 2017; Amorín et al. 2017; Nanayakkara et al. 2019). The EWs of He II, O III]λ1666, and C III] from these works are larger than those measured in the average spectra of our LAEs, with the exception of some of the MUSE sources at z ≳ 3 studied by Patrício et al. (2016) and Nanayakkara et al. (2019), some of the stacks from Nakajima et al. (2018b), and some of the VANDELS He II emitters from Saxena et al. (2020). The EWs from Nakajima et al. (2018b) are computed from average spectra of a z ≈ 3 LAE population whose median UV luminosity is about two orders of magnitude brighter than ours. Most of the 2.4   z   3.5 UV-selected, low-luminosity galaxies from Amorín et al. (2017) have higher He II and C III] EWs, most likely because of their higher SFR (see Fig. 2). The line measurements from Saxena et al. (2020) are performed on the stacked spectra of UV continuum, bright He II emitters (−19   MUV   −22). The strong line emitters from Erb et al. (2010) and Berg et al. (2018) are brighter and more massive than the median value of our LAEs. The lensed galaxies from Vanzella et al. (2016, 2017) are among the least massive, most metal-poor, young, and faintest systems observed at z ∼ 3. With MUV >  −16, the source ID14 from Vanzella et al. (2017) is roughly one order of magnitude fainter that the faintest LAE in our sample, while the source ID11 from Vanzella et al. (2016) has an remarkable blue UV slope (β = −2.95). Recently, Du et al. (2020) measured C III] EWs in z ∼ 2 analogs of galaxies in the reionization era, obtaining values from 13.2 Å down to 1 Å, reaching values as low as those of our LAEs. The authors also found differences in the C III] EW depending on the selection criteria, with higher values of C III] EW for emission lines rather than for continuum-selected galaxies. Mainali et al. (2020) detected two targets with C III] emission reaching EW ≈ 17 − 21 Å in the spectra of z ∼ 2 galaxies selected for their strong rest-optical line ([O III]λ5007+Hβ) EWs. These values are above the ones of our LAEs and are similar to those observed at z >  6 (e.g., Stark et al. 2015b, 2017; Hutchison et al. 2019).
Citation: Nakajima et al. 2018b; Nakajima et al. (2018b); Nakajima et al. (2018b)
Citation function: Background; Background; Similarities; Similarities; Differences; Differences
Here we shall set the potential 





 and make the change of variable 





; hence, the vortical parts of the footpoint equations of motion are cast into a Hamiltonian form. The Hamiltonian describing the dynamics of magnetic footpoints driven by the vortical component of the turbulence on the photosphere is
50





The footpoint Hamiltonian 





 has the physical dimension of a frequency (s−1). From Figure 2 in Rincon et al. (2017), we further notice that the vortical component of the photospheric surface velocity field has a structure reminiscent of 2D Euler turbulence, which admits a vortex point representation (Eyink & Sreenivasan 2006) also on the sphere (Pavlov et al. 2001). There are three important parameters that can be extracted from the vortical component of photospheric turbulence. These parameters are the typical amplitude of the velocity fluctuations 





, the correlation length 





, and the correlation time 





 of the turbulence, where 





 is the rms amplitude of the fluctuating stream function Φ and 





 is the angular size of the turbulent eddies on the source surface. These three independent parameters 





, 





, and 





 can all be derived from a more fundamental quantity, which is the covariance function 





, i.e., the autocorrelation function of the turbulent stream function (Creasey & Lang 2018). The covariance function is sufficient to characterize the vortical component of the surface flows provided that the stream function Φ does not significantly deviate from a stationary isotropic Gaussian random field on the sphere. Hamilton equations with the Hamiltonian (50) determine a genuine problem of chaotic mixing (Ott 1993; Morrison 1998) on the sphere. We emphasize that mixing of footpoints on the photosphere relies on the time dependence of the Hamiltonian. In fact, the footpoint dynamics is regular (quasi-periodic) when the Hamiltonian is a constant of time, independently of the number of modes involved in its decomposition. The reason is that the dynamics generated by 1 degree of freedom Hamiltonians, here for the pair of canonically conjugate variables μ and ϕ, is regular. A simple mathematical proof of this statement is given by Morrison (1998). In one version of the protocol used by Giacalone & Jokipii (2004) to model the magnetic footpoint dynamics, the time dependence of the flow results from the (random) phase of the (complex) toroidal coefficients being renewed at random times, following the work of Simon et al. (1995).
Citation: Eyink & Sreenivasan 2006
Citation function: Uses; Uses
We assume that gas and dust are well mixed along each line of sight in which case the gas surface density (Σgas) is proportional to the dust surface density (Σdust) and the proportionality factor is the gas-to-dust mass ratio (rgd). The dust masses derived by SED fitting methods vary significantly and systematically (e.g. Galliano et al. 2011; Gordon et al. 2014) depending on the applied method and assumed dust optical properties. It is therefore important to use a value of rgd consistent with the used SED fitting methodology, i.e. the value of rgd that gives the correct gas mass given the derived dust mass. We have calibrated the rgd using the available gas tracers on a large scale. Within a 500 arcsec radius centred on NGC 346, we find an atomic gas mass of 2.7×106 M⊙ using the Parkes+ATCA H i map (Kim et al. 2003) and a spin temperature of 60 K (Bernard et al. 2008). The molecular gas mass determined from CO(J = 1–0) within the same aperture is 7.4×104 M⊙ when using a CO-to-H2 column density conversion factor (XCO) of 1021 [cm−2 (K km s−1)−1] (Bolatto, Wolfire & Leroy 2013). If we use an XCO of 1022 [cm−2 (K km s−1)−1], we find a molecular gas mass of 7.5×105 M⊙. Such an elevated XCO may be more appropriate given the harsh radiation field that is prevalent in the region, which may induce enhanced photodissociation of the CO molecules. Thus, we derive a total gas mass in the range of 2.7×106–3.4×106 M⊙. The dust mass derived using the dust continuum photometry in this aperture and applying the same SED fitting routine we use for the pixel-by-pixel analysis yields 2.7×103 M⊙. The derived rgd ranges from ∼1000 to ∼1250. In the following, we use a fiducial value for rgd of 1250, which is much higher than the Galactic value (100; Draine & Li 2007) due to the lower metallicity of the SMC. Our value for rgd of 1250 is lower than the total SMC integrated value of 1740 from Gordon et al. (2014). The global value includes a significant amount of atomic gas at low column density in the outskirts of the SMC which is not detected in dust emission. This drives up the global rgd.
Citation: Kim et al. 2003
Citation function: Uses
Here, we report on molecular gas observations of NGC 0708, the BCG in the low-mass galaxy cluster Abell 262, itself part of the Perseus–Pisces galaxy supercluster. NGC 0708 lies 58.3 ± 5.4 Mpc away (estimated using infrared surface brightness fluctuations; Jensen et al. 2003). It is a giant elliptical galaxy with a weak dust lane (Ebneter & Balick 1985; Wegner et al. 1996) and an effective radius of 33 arcsec ($\approx \, 9.3$ kpc; Wegner et al. 2012). See Fig. 1 for an HST image of NGC 0708. Abell 262 was identified as having an X-ray emitting ICM by Jones & Forman (1984), and Stewart et al. (1984) measured the cooling time to be 1.3 × 109 yr, smaller than the age of the Universe so that the cluster is expected to form a cooling flow. The 20-cm observations of Parma et al. (1986) revealed a double-lobed, ‘S’-shaped jet and led to the classification of NGC 0708 as a weak Fanaroff–Riley Class I radio source (Blanton et al. 2004). The top panel of Fig. 1 also has 330 MHz continuum observations from Clarke et al. (2009) overlaid (blue contours) to show the shape and orientation of the large-scale jet. Analysis of Chandra observations revealed a hole or bubble within the ICM, cospatial with the eastern lobe of the jet (Blanton et al. 2004). Clarke et al. (2009) found additional 3–6 kpc radius cavities at differing position angles within the X-ray gas, and at a range of radial distances from the BCG (8–29 kpc), indicating multiple episodes of AGN activity from a precessing SMBH jet. They concluded that the total AGN emission should be capable of counteracting the cooling flow over several outbursts. Using their multifrequency observations of NGC 0708, Clarke et al. (2009) also calculated the radio spectral index (α) from 235 to 610 MHz, finding the spectrum to be flat in the core (α = −0.5), typical of new particles in a jet. They also estimated a lower limit on average outburst repetition time-scales in Abell 262 to be τrep ≥ 28 Myr.
Citation: Blanton et al. 2004
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C18O (J = 2–1) observations revealed a rotational structure in the north–south direction at the center of IRAS 4C, but whether the rotation is Keplerian is not known due to the low signal-to-noise ratio (S/N) (Tobin et al. 2015). Spitzer IRAC observations show an outflow cavity structure highlighted by scattered light and shocked emission to the east side of the central source, with the west side being much fainter (Figure 19 of Tobin et al. 2015). The east side is therefore inferred to be the blueshifted side, because the blueshifted outflow cavity tends to be brighter than the redshifted cavity in NIR and MIR as it is less extincted. Despite the outflow cavity structure seen in infrared, there was no clear evidence of a molecular outflow from mm or submm observations. Previous 12CO observations either reported no detection of outflow emission toward this source (Plunkett et al. 2013; Tobin et al. 2015) or only weak compact blueshifted emission towards east of the continuum source (Stephens et al. 2018). One possible explanation for the weak 12CO outflow emission is that the outflow lies close to the plane of sky, so that the low-velocity outflow emission is easily mixed with the emission of the ambient gas, especially for abundant species like 12CO. Indeed, the inclination of the source is estimated to be nearly edge-on (Tobin et al. 2015) to about 25° between the disk plane and the line of sight (Segura-Cox et al. 2016). On the other hand, the 13CO (J = 2–1) emission reveal a compact (≲2″) structure with the blueshifted emission slightly offset to the east of the redshifted emission, which was explained as a slow outflow (Koumpia et al. 2016). Here, we report that the outflow cavity structure is clearly detected in the CCH and CS emissions, with kinematics consistent with rotation with respect to the outflow axis. This allows us to measure the angular momentum distribution in the outflow, and further constrain its launching radii and launching mechanism.
Citation: Plunkett et al. 2013
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