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RCW 57A (also known as NGC 3576, G291.27–0.70, or IRAS 11097–6102) is a H ii region associated with a filament and bipolar bubble, and is located at a distance of 2.4–2.8 kpc (Persi et al. 1994; de Pree et al. 1999). We adopt 2.4 kpc, which is within uncertainties of both the kinematic and spectroscopic determinations (see Persi et al. 1994). Figure 1 depicts the overall morphology of RCW 57A. It contains optically bright nebulosity with several dark globules and luminous arcs (Persi et al. 1994). It is one of the massive star-forming regions in the southern sky, hosting a H ii region (cyan contour) embedded in a filament (white contour) from which a widely extended bipolar bubble (yellow contours) is emerging. A deeply embedded near-IR cluster, consisting of more than 130 young stellar objects (YSOs), is associated with this region (Persi et al. 1994). The observed ratios of the infrared fine-structure ionic lines (Ne ii, Ar iii, and S iv; Lacy et al. 1982) indicate that at least eight O7.5V stars are necessary to account for the ionization of the region. However, even these stars may not be sufficient to account for the Lyα ionizing photons inferred from radio data (Figuerêdo et al. 2002; Barbosa et al. 2003; Townsley 2009). Based on the newly discovered cluster of stars, using X-ray data, Townsley et al. (2014) suggested that an additional cluster of OB stars that were not known before might be deeply embedded. This cluster is located slightly southwest of the center of the near-IR cluster. The 10 μm map (cf., Frogel & Persson 1974) reveals the presence of five infrared sources (IRS; black squares in Figure 1) near the center of the H ii region. These, together with water and methanol maser sources (green crosses in Figure 1) distributed along the filament, are indicative of active ongoing star formation in RCW 57A (Frogel & Persson 1974; Caswell 2004; Purcell et al. 2009). Therefore, RCW 57A is an ideal target to investigate the morphological links among filaments, bipolar bubbles, and B-fields so as to understand the star formation history.
Citation: Lacy et al. 1982
Citation function: Uses; Uses
The characterization of the interior of observed exoplanets is one of the main goals in current exoplanetary science. With the large number of newly discovered exoplanets expected in the next 10 years by ground-based surveys such as the Wide Angle Search for Planets (WASP; Pollacco et al. 2006), the Next-Generation Transit Survey (NGTS; Wheatley et al. 2017), and the Hungarian Automated Telescope Network/Hungarian Automated Telescope Network-South (HATNet/HATSouth; Hartman et al. 2004; Bakos et al. 2013), as well as the ongoing Transiting Exoplanet Survey Satellite (TESS) space survey (Ricker et al. 2014) and the upcoming PLAnetary Transits and Oscillations (PLATO) mission (Rauer et al. 2014), a rapid characterization scheme of the interior structure of these planets will become increasingly necessary to further our understanding of planetary populations. The vast majority of the confirmed exoplanets has been identified either through transits or radial velocities surveys. Planets identified with both techniques are characterized by their mass and radius, which, combined, provide a first indication of the bulk composition through comparison with theoretical mass–radius curves (e.g., Valencia et al. 2006; Sotin et al. 2007; Zeng & Sasselov 2013). A common approach to the interior characterization of exoplanets is the use of numerical models to compute interior structures that comply with the measured mass and radius of the planet (e.g., Fortney et al. 2007; Sotin et al. 2007; Valencia et al. 2007; Wagner et al. 2011; Zeng & Sasselov 2013; Unterborn & Panero 2019). As this is an inverse problem, it requires the calculation of a large number of interior models to obtain an overview over possible interior structures (Rogers & Seager 2010a; Brugger et al. 2017; Dorn et al. 2017). If other observables are used in addition to mass and radius, the number of samples needed for an accurate inference of possible interior structures increases drastically, due to the increase in dimensionality (e.g., James et al. 2013). Thus, the inference can quickly become computationally expensive. Moreover, with only mass and radius, possible solutions tend to be highly degenerate, with multiple, qualitatively different interior compositions that can match the observations equally well (e.g., Rogers & Seager 2010a, 2010b).
Citation: Ricker et al. 2014
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The object J004415.00 required our special consideration. We notice large variations of spectra at the hydrogen series limits. We designate the limits of the Balmer, Paschen, Brackett and Pfund series in Fig. 4 for J004415.00 with vertical arrows. A jump at the Brackett series limit in the IR spectrum can be clearly seen. The Balmer and Paschen jumps can also be noticed by photometric brightness changes. These inverse jumps at the hydrogen series limits attest to the noticeable contribution of free–free (f–f) and free–bound (f–b) emissions to the spectrum. This confirms the presence of an ionized circumstellar envelope where the f–f and f–b emissions originate, which is typical for B[e]SGs (Zickgraf et al. 1985, 1986, 1989). To take the contribution of f–f and f–b radiations into account and estimate the corresponding spectra, we used the chianti package (Dere et al. 1997; Landi et al. 2013). We consider the case of isothermal pure hydrogen plasma at a temperature of Te = 10 000 K (Lamers et al. 1998). From the spectrum of the star we assume that its effective temperature is Tstar =15 000 K and match the model spectrum to the observations with the following parameters: Te = 10 000 K, EM = 1.37 × 1039 cm−5 (emission measure), Tstar = 15 000 K, Rstar = 43R⊙, AV = 1.1. In order to describe the IR excess of the SED we also add hot dust emission to the model spectrum by assuming it to be effective blackbody radiation with a temperature of about 1000 K (the best fit obtained with Tdust = 1050 K), which is typical for B[e]SG hot circumstellar dust (Lamers et al. 1998). The result of the SED fitting is shown in Fig. 4: the dashed line shows the blackbody radiation, the dotted line designates the contribution of free–free and free–bound emissions, the dash-dotted line demonstrates the dust emission contribution, the solid line indicates the total model spectrum. All spectra shown include the interstellar extinction (Fitzpatrick 1999, assuming RV = 3.1) with AV = 1.1 obtained as the best-fitting parameter. Since the object J004415.00 does not show considerable photometric variability, we incorporate some more published photometric data points for fitting (see Table 3) in addition to the BTA data. For better illustration of the matching between the model and observed spectra, the synthetic photometry of the total model continuum multiplied by the normalized spectrum is shown with filled circles. Note that since the parameters are obtained in the isothermal plasma approximation with fixed temperatures of the star and circumstellar plasma, we report the other parameters derived from them as estimates and do not assess their uncertainties. Nevertheless, our estimates of the model parameters are good enough and allow us to constrain the dust extinction and luminosity of the object.
Citation: Dere et al. 1997
Citation function: Uses
Cluster radio relics are usually found in the outskirts of merging galaxy clusters. They exhibit elongated morphologies and high degrees of polarisation above 1 GHz (up to 70%, Ensslin et al. 1998; Bonafede et al. 2014; Loi et al. 2019; de Gasperin et al. 2022). The resolved spectral index in radio relics shows a gradient: it steepens towards the cluster centre and flattens towards the outskirts. Their size can reach up to ∼2 Mpc, and high-resolution observations have revealed filamentary structures within relics themselves (Di Gennaro et al. 2018; Rajpurohit et al. 2020, 2022a,b; de Gasperin et al. 2022). The Largest Linear Sizes (LLS) and radio powers of relics are correlated, as well as the integrated spectral index and the radio power (van Weeren et al. 2009b; Bonafede et al. 2012; de Gasperin et al. 2014). Relics trace ICM shock waves with relatively low (M   3) Mach numbers (Finoguenov et al. 2010; Akamatsu et al. 2013; Shimwell et al. 2015; Botteon et al. 2016). The acceleration of electrons is believed to proceed via diffusive shock acceleration (DSA) in the ICM (Ensslin et al. 1998; Roettiger et al. 1999), in which particles scatter back and forth across the shock front gaining energy at every crossing. Nevertheless, this mechanism has been shown to be rather inefficient in accelerating electrons from the thermal pool (Vazza & Brüggen 2014; Vazza et al. 2016; Botteon et al. 2020a; Brüggen & Vazza 2020; see Brunetti & Jones 2014 for a review). Recently, it has been suggested that seed electrons could originate from the tails and lobes (driven by AGN outflows) of cluster radio galaxies (Bonafede et al. 2014; van Weeren et al. 2017; Stuardi et al. 2019), which alleviates the requirements of high acceleration efficiencies at cluster shocks (e.g., Markevitch et al. 2005; Kang et al. 2012, 2017; Botteon et al. 2016; Eckert et al. 2016). In some cases, double relics have been detected on opposite sides of the cluster centre (e.g., Rottgering et al. 1997; van Weeren et al. 2010, 2012b; Bonafede et al. 2012; de Gasperin et al. 2015a). In these clusters it is possible to constrain the merger history, providing important information about the formation processes of relics.
Citation: Finoguenov et al. 2010
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Citation: Bonafede et al. 2012
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Citation: Bonafede et al. 2012
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Citation: van Weeren et al. 2009b
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For comparisons with the light curve YORP constraints, the YORP effect acting on Cuyo could be predicted by computing the total recoil forces and torques from reflected and thermally emitted photons from the asteroid surface using the ATPM. These calculations were made for both a smooth and rough surface, and were averaged over both the asteroid rotation and elliptical orbit (see Rozitis & Green 2012, 2013, for methodology). As demonstrated in Rozitis & Green (2012), the inclusion of rough-surface thermal-infrared beaming effects in the YORP predictions tends to dampen the YORP rotational acceleration on average but can add uncertainties of up to several tens of per cent if the roughness was varied across the surface. Since the light curve inversion produced convex shape models only, then shadowing and self-heating effects inside global-scale concavities (see Rozitis & Green 2013) were not possible to model. However, a study of non-convex shape models for fast two to four hour rotators in Rozitis & Green (2013) indicated that such asteroids have rather minimal levels of global-scale concavities, and the ~ 2.7 h rotation period of Cuyo implies that its shape could be similar. Furthermore, the Tangential-YORP effect, that is, a predicted rotational acceleration caused by temperature asymmetries within exposed rocks and boulders on the surface of an asteroid (Golubov & Krugly 2012), was also not included in the ATPM predictions. However, the very low thermal inertia value measured for Cuyo implies the absence of rocks and boulders on its surface of the quantity and size that are necessary to induce a significant Tangential-YORP component. As Cuyo is likely to be an S-type rubble-pile asteroid, a bulk density equivalent to that measured for the S-type rubble-pile asteroid (25143) Itokawa (Abe et al. 2006) of 2 g cm−3 was assumed for the YORP computations. Using the thermo-physical properties derived earlier, the ATPM predicts YORP rotational acceleration of (−6.39 ± 0.96) × 10−10 rad day−2 for the nominal shape model. The uncertainty given here corresponds to the standard deviation of results when the degree of surface roughness israndomly varied across the surface of Cuyo (see Lowry et al. 2014, for details of the Monte Carlo methodology used). These values lie well within the light curve rotational acceleration constraints determined previously.
Citation: Rozitis & Green (2013)
Citation function: Similarities; Similarities
Citation: Rozitis & Green 2013
Citation function: Differences; Differences
The understanding of proton acceleration at the Sun in large solar energetic particle (SEP) events has oscillated between flare and shock pictures (Cliver 2009b; Reames 2015). The earliest picture following the discovery of ground-level events (GLEs; major SEP events requiring >500 MeV protons) by Forbush (1946) was that protons were accelerated in flares, the clear choice in the absence of other observations. Subsequently, Wild et al. (1963) conjectured, mainly on the basis of radio observations, that large SEP events required coronal shock waves as manifested by type II solar radio bursts. Smaller electron-dominated SEP events were linked to metric type III bursts. Early observational support for this view on the SEP side was provided by Lin (1970). Through the work of Švestka & Fritzová-Švestková (1974), Kahler et al. (1978, 1984), Cliver et al. (1982, 1983a, 1983b), Cane & Stone (1984), Klecker et al. (1984), Mason et al. (1984, 1986), Meyer (1985), Reames et al. (1985, 1994, 1996), Cane et al. (1986, 1988), Luhn et al. (1987), Reames (1990, 1999), Kahler (1992, 1994), Gosling (1993), and others, involving various comparisons of SEP events with flare electromagnetic emissions and coronal mass ejections (CMEs), as well as considerations of SEP composition, charge states, and the longitude distribution of SEP-associated flares, the two-class picture of SEP acceleration presciently proposed by Wild et al. (1963) became established. The new consensus view was almost immediately challenged by observations of the first large (“gradual”; Reames 1993) proton events observed by the Advanced Composition Explorer (ACE). Mazur et al. (1999), Cohen et al. (1999), Mason et al. (1999a), and Mason et al. (1999b) reported that large SEP events, including GLEs, recorded by ACE and SAMPEX in 1997 and 1998 had elemental composition and charge states at >10 MeV/nuc that were similar to those found in small (“impulsive”) SEP events (e.g., Mason et al. 1986; Luhn et al. 1987) at lower energies. Subsequently, Cane et al. (2002, 2003, 2006) presented evidence based on low-frequency radio observations, SEP composition data, and flare location to argue for the presence of a flare-accelerated high-energy (>25 MeV) proton component in large SEP events to augment that produced by coronal/interplanetary shock waves driven by CMEs. The relative importance of flare and shock components was left as an open question.
Citation: Cliver et al.; 1983a
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The core of a neutron star is composed primarily of neutrons, with ∼5%–10% of the mass in protons; for the electrically neutral medium, the number density of electrons is equal to that of the protons. At the supranuclear densities of the outer core, the Fermi energy for protons and neutrons is well above the typical temperature of a mature neutron star, and both the neutrons and protons are expected to condense into Bardeen–Cooper–Schrieffer superfluids, with 3PF2 and 1S0 Cooper pairing, respectively (Migdal 1959; Baym et al. 1969). To support rotation, the neutron superfluid forms an array of quantized vortices, filaments of microscopic cross section, each carrying one quantum of circulation. The superconductivity of the protons is predicted to be type II, and the magnetic field is supported by an array of quantized flux tubes, each carrying one quantum of magnetic flux. Fermi-liquid interactions between the two condensates results in a nondissipative coupling between the mass currents of the two species (Andreev & Bashkin 1975; Chamel & Haensel 2006), so the neutron vortices are magnetized by entrained proton currents (Alpar et al. 1984a). Electron scattering from magnetized vortices and flux tubes produces dissipative and nondissipative forces on the vortices and flux tubes. The magnetic interaction at junctions between magnetized neutron vortices and flux tubes is energetic enough to produce pinning, wherein the neutron vortices pin to the dense array of flux tubes in the outer core (Srinivasan et al. 1990; Jones 1991; Chau et al. 1992; Ruderman et al. 1998; Link 2012b), similar to the predicted pinning of the vortices to the nuclear lattice of the crust (Anderson & Itoh 1975; Alpar 1977; Epstein & Baym 1988; Donati & Pizzochero 2006; Avogadro et al. 2007; Link 2009). Thermal fluctuations stochastically excite vortex motion, causing the neutron vortices to slip with respect to the flux tubes (Ding et al. 1993; Sidery & Alpar 2009; Link 2014).
Citation: Baym et al. 1969
Citation function: Uses
We consider a radio galaxy of SMBH mass M = 109
M
9
M
⊙ with a mass accretion rate of 



Ṁ=ṁLEdd/c2≃1.4×1022M9ṁ−4gcm−2

, where c is the speed of light and L
Edd is the Eddington luminosity. The gravitational radius of the BH is r

g
 = GM/c
2 ≃ 1.5 × 1014
M
9 cm. We consider that the accretion flow is in the MAD state, and then the magnetic field strength around the SMBH is estimated to be 



Bmad=ṀcΦmad2/(4π2rg2)≃1.1×103M9−1/2ṁ−41/2Φmad,1.7G

 (e.g., Yuan & Narayan 2014), where Φmad ≈ 50Φmad,1.7 is the saturated magnetic flux (Tchekhovskoy et al. 2011; Narayan et al. 2012; McKinney et al. 2012; White et al. 2019). The high-resolution GRMHD simulation with a BH spin parameter a = 0.9375 suggests that magnetic reconnection occurs at a distance of r
rec ∼ 2r

g
 (Ripperda et al. 2022). The value of r
rec could depend on a or other parameters, but we fix r
rec = 2r

g
 throughout this paper for simplicity. We estimate the reconnecting magnetic field strength to be (see Appendix A)
1



Brec≈2Bmadrrecrg−2≃3.9×102M9−1/2ṁ−41/2Φrec,1.2G,

where 



Φrec=2Φmad(rrec/rg)−2

 is the effective magnetic flux at the reconnection region. The magnetosphere will be formed around the SMBH. The minimum number density of the magnetosphere that can maintain the electric current for the BZ process is (Goldreich & Julian 1969; Levinson & Cerutti 2018)
2



nGJ=BrecΩF2πec≈Brec8πerg≃2.2×10−4M9−3/2ṁ−41/2Φrec,1.2cm−3,

where Ω
F
 ≈ ac/(4r

g
) is the field line angular velocity (Tchekhovskoy et al. 2010; Nathanail & Contopoulos 2014; Ogihara et al. 2021; Camilloni et al. 2022) and we assume the BH spin parameter as a ∼ 1. For the magnetosphere, which consists of e
+
e
− pair plasma with the density n
GJ, the magnetization parameter is
3



σB,GJ=Brec24πnGJmec2≈6.8×1013M91/2ṁ−41/2Φrec,1.2.

This value should be regarded as an upper limit, because the number density of the magnetosphere can be higher than n
GJ. Various mechanisms of particle injection into the BH magnetosphere have been proposed (see Appendix B), which can lead to multiplicity of κ
± ≡ n/n
GJ ∼ 1 − 103. This results in the magnetization parameter of σ

B
 ≳ 1010.
Citation: Narayan et al. 2012
Citation function: Uses