adsabs/FOCAL huggingface dataset:
PYU+ formed from HC3N and H2NCO+, having at least 88 kJ mol−1 of internal energy, can redissociate or be stabilized by radiative relaxation. Collisional stabilization in the ISM is not effective due to an extremely low gas density. To compare rates of redissociation and radiative relaxation, the microcanonical rate constant for PYU+ → HC3N + H2NCO+ was calculated using the statistical Rice–Ramsperger–Kassel–Marcus (RRKM) theory (Marcus & Rice 1951). The method has been described elsewhere (Baer & Hase 1996; Yim & Choe 2011; Choe et al. 2017; Choe & Kim 2019). At 10–50 K, energies in the range of 0.2–1 kJ mol−1 (5/2RT) of thermal translational and rotational energies of HC3N and H2NCO+ are transferred to an vibrational internal energy of PYU+, which are added to 88 (87.8, more precisely) kJ mol−1 of the energy of PYU+ for redissociation. RRKM rate constants calculated at this energy range were (3–7) × 103 s−1 (Table 1). On the other hand, rate constants for typical radiative relaxations of vibrationally excited ions have been found to be lie within the range 101–103 s−1 (Herbst 1985; Dunbar et al. 1996; Smith et al. 2001; Wakelam et al. 2010). This indicates that a considerable number of the excited PYU+ ions would be stabilized by comparing rate constants. At lower temperatures, stabilization would be more effective because rate constants for dissociation and radiative relaxation are expected to show positive and negative dependences on temperature, respectively (Wakelam et al. 2010). The stabilized PYU+ ions can undergo a further association reaction with NH3, H2O, or CH3OH. We will call this case A. These association reactions with NH3 and H2O, which finally lead to CyH+ and UrH+, respectively, can occur without activation barriers. The latter reaction is barrierless when assisted by a H2O molecule. On the other hand, the reaction of stabilized PYU+ ion with CH3OH to form ThH+ can hardly occur thermally at low temperatures due to its overall activation energy of 13 kJ mol−1.
Citation: Choe et al. 2017
Citation function: Background
The availability of two rotational transitions with different upper level energies allows us to constrain the rotational temperature of the propargyl radical in TMC-1. We built a rotation diagram using the velocity-integrated intensities given in Table 1, and we derive a rotational temperature of 9.9 ± 1.5 K (see Fig. 2). We therefore confirm the assumption made by Agúndez et al. (2021a) that the rotational levels of CH2CCH are thermalized at the gas kinetic temperature of TMC-1, ~10 K (Fehér et al. 2016). This fact is expected based on the low dipole moment of CH2CCH (0.150 D; Küpper et al. 2002), which implies low critical densities, probably a few 102 cm−3 (i.e., well below the volume density of H2 in TMC-1, afew 104 cm−3; Pratap et al. 1997; Cordiner et al. 2013). The column density derived from the rotation diagram for ortho CH2CCH is (8.2 ± 1.7) × 1013 cm−2. A more precise determination of the column density can be obtained by fitting the observed spectra with synthetic spectra calculated under local thermodynamic equilibrium. For this calculation we adopted a rotational temperature of 9.9 K, as derived from the rotation diagram, a full width at half maximum of 0.72 km s−1 for the 20,2–10,1 lines and 0.57 km s−1 for the 50,5–40,4 lines, which are the arithmetic mean of the values derived for the hyperfine components of each transition (see Table 1), and assumed that the emission is distributed in the sky as a circle with a radius of 40″, as observed for various hydrocarbons in TMC-1 (Fossé et al. 2001). The observed spectra at 37.5 GHz and 93.6 GHz are well reproduced when adopting a column density of 7.5 × 1013 cm−2 (see Fig. 1). Assuming an ortho-to-para ratio of three, the column density of CH2CCH (including ortho and para) in TMC-1 is (1.0 ± 0.2) × 1014 cm−2, which is slightly higher than the value derived previously by Agúndez et al. (2021a). The column density of the closed-shell counterpart CH3CCH in TMC-1 is (1.1–1.3) × 1014 cm−2 (Gratier et al. 2016; Cabezas et al. 2021). Therefore, in this study we confirm that the propargyl radical is thermalized to the gas kinetic temperature of TMC-1 and revise the abundance ratio CH2CCH/CH3CCH to nearly one.
Citation: Küpper et al. 2002
Citation function: Similarities; Uses
Over the last 15 yr, our understanding of how galaxies form and evolve has improved substantially. The combination of technological and theoretical progress has brought this field into a new era: advances in observational techniques (e.g. multi-object spectroscopy, efficient near-infrared CCDs) have enabled multiwavelength observations of large samples of galaxies out to the highest redshifts, while the steady rise of computational power and refinement of numerical techniques have fostered new approaches (e.g. semi-analytic models, hydro-dynamic simulations) to model the formation and evolution of galaxies. This progress has led to a general consensus about the main physical ingredients required to describe the evolution of the galaxy population (e.g. Gonzalez-Perez et al. 2014; Lu et al. 2014; Vogelsberger et al. 2014; Henriques et al. 2015; Schaye et al. 2015): collapse and hierarchical growth of dark matter haloes; accretion of baryons on to these haloes; conversion of baryons into stars; feedback of massive stars and active galactic nuclei (AGN) on star formation; supernova- and AGN-driven outflows of metal-enriched gas; infall of both pristine and metal-enriched gas on to galaxies. The large-scale environment can also affect galaxy properties, in particular, by providing quenching mechanisms (e.g. tidal or ram-pressure stripping, strangulation; e.g. Lagos et al. 2014; Rafieferantsoa et al. 2015), and through its influence on the merger rate (e.g. Lackner et al. 2012; Rafieferantsoa et al. 2015) and galactic spins (e.g. Hahn, Teyssier & Carollo 2010; Codis et al. 2012). Although these different ingredients are present in many galaxy formation models, we still lack a detailed quantification of their respective roles in shaping the properties of galaxies. This is because of the complexity inherent in galaxy physics, which combines gravity, radiation hydro-dynamics, magnetic fields and high-energy physics, acting on scales from less than a pc (e.g. for the formation of proto-stellar cores) to over a Mpc (e.g. for environmental effects). For this reason, ‘first-principles’ simulations of galaxy formation remain far beyond the reach of current computational capabilities. Instead, small-scale baryonic physics is generally subsumed into sub-grid prescriptions, which vary from model to model (e.g. Scannapieco et al. 2012; Haas et al. 2013a,b; Vogelsberger et al. 2013; Torrey et al. 2014; Crain et al. 2015). The appropriateness of such prescriptions, and hence, our ability to understand galaxy formation, must be assessed by comparing simulated and observed galaxy properties.
Citation: Gonzalez-Perez et al. 2014
Citation function: Background; Background
The hot emission line of Fe xxi 1354.09 Å and the cool emission line of Si iv 1402.77 Å have been used in many spectroscopic studies to investigate chromospheric evaporation (e.g., Tian et al. 2014, 2015; Li et al. 2015b, 2017a, 2017b; Brosius et al. 2016; Zhang et al. 2016a, 2016b). It is widely accepted that the forbidden line of Fe xxi 1354.09 Å is a hot (log T ∼ 7.05) and broad emission line during solar flares (Doschek et al. 1975; Cheng et al. 1979; Mason et al. 1986; Innes et al. 2003a, 2003b). Meanwhile, IRIS spectroscopic observations show that Fe xxi 1354.09 Å is always blended with a number of cool and narrow emission lines, which are from neutral or singly ionized species. Those blended emission lines can be easily detected at the position of the flare ribbon, including known and unknown emission lines, such as the C i line at 1354.29 Å, the Fe ii lines at 1353.02 Å, 1354.01 Å, and 1354.75 Å, the Si ii lines at 1352.64 Å and 1353.72 Å, and the unidentified lines at 1353.32 Å and 1353.39 Å (e.g., Li et al. 2015a, 2016a; Polito et al. 2015, 2016; Tian et al. 2015, 2016; Young et al. 2015; Tian 2017). In order to extract the hot line of Fe xxi 1354.09 Å and the cool line of C i 1354.29 Å (log T ∼ 4.0; Huang et al. 2014), we apply a multi-Gaussian function superimposed on a linear background to fit the IRIS spectrum at the “O i” window (e.g., Li et al. 2015a, 2016a), which has been pre-processed (i.e., IRIS spectral image deformation, bad pixel despiking and wavelength calibration) with the standard routines in Solar Soft Ware (SSW; Freeland et al. 2000). In short, the line positions and widths of these blended emission lines are fixed or constrained, and their peak intensities are tied to isolated emission lines from similar species. More details can be found in our previous papers (Li et al. 2015a, 2016a). On the other hand, the cool line of Si iv 1402.77 Å (log T ∼ 4.8) at the “Si iv” window is relatively isolated, and it can be well fitted with a single-Gaussian function superimposed on a linear background (Li et al. 2014, 2017a). Using the relatively strong neutral lines (i.e., “O i” 1355.60 Å and “S i” 1401.51 Å), we also perform an absolute wavelength calibration for the spectra at the “O i” and “Si iv” windows, respectively (Tian et al. 2015; Tian 2017). Finally, the Doppler velocities of Fe xxi 1354.09 Å, C i 1354.29 Å, and Si iv 1402.77 Å are determined by fitting line centers removed from their rest wavelengths (Cheng & Ding 2016b; Guo et al. 2017; Li et al. 2017a). As the hot Fe xxi line is absent in the non-flaring spectrum, the rest wavelength for the Fe xxi line (i.e., 1354.09 Å) is determined by averaging the line centers of the Fe XXI profiles which were used in the previous IRIS observations (Brosius & Daw 2015; Polito et al. 2015, 2016; Sadykov et al. 2015; Tian et al. 2015; Young et al. 2015; Brosius et al. 2016; Lee et al. 2017), while the rest wavelengths for the C i and Si iv lines, i.e., 1354.29 Å and 1402.77 Å, respectively, are determined from their quiet-Sun spectra (Li et al. 2014, 2015a).
Citation: Li et al. 2015a
Citation function: Uses; Uses; Uses
In the 2D density distribution (Fig. 5, upper panel) we can distinguish only the primary structure. Taking into account the measured mass according to the fitted shear profile for this component, we compute the corresponding projected density profile, Σ(r), assuming a NFW distribution according to the equations given by Wright & Brainerd (2000). All of the obtained values are below the 3σ detection level (250 h70 M⊙ pc−2) and only the inner region, corresponding to a radius of 60 kpc would present a density larger than the 2σ detection level (180 h70 M⊙ pc−2). Thus, we consider that the mass of this structure is near the detection threshold of the 2D density contrast distribution. Also, we do not detect significant X-ray emission in the secondary structure region above the threshold adopted to build the brightness contours. In order to establish if this is due to the observing detection limit, we consider the lowest detected flux of RASS-based catalogs according to Piffaretti et al. (2011), which contains the lowest X-ray emission clusters identified using this data. This corresponds to a flux Flim = 1.5 × 10−12 erg s−1 cm−2, in the 0.1–2.4 keV band of the SGP catalog (Cruddace et al. 2002). Considering this flux and the redshift of the secondary structure, we obtain a limiting luminosity of 1.5 × 1042 


$h^{-2}_{70}$



h

70

−2



 erg s−1. Taking into account the obtained lensing mass, we compute the expected X-ray luminosity for the system according to Leauthaud et al. (2007) M_ LX relation, obtaining (3.7 ± 2.4) × 1043 


$h^{-2}_{70}$



h

70

−2



 erg s−1, only 1.5σ above the detection limit. For this calculation we use the derived SIS mass since it has the lowest χ2 value and errors correspond to the propagation of the M_LX parameter errors. Therefore, according to the lensing estimated mass the secondary structure would be roughly at the detection limit of the X-ray observation. We considered using more sensitive X-ray data, therefore we checked in the available databases of Chandra and XMM; however this component is not present in the field-of-view of these surveys. Considering the X-ray luminosity emission upper limit (6.1 × 1043 


$h^{-2}_{70}$



h

70

−2



 erg s−1) the emission is 2.5σ above the detection limit given by Flim (Piffaretti et al. 2011). Therefore, it is important to highlight that a low X-ray emissivity could be explained by a low density of the intracluster gas, which might be produced by a past interaction between the structures.
Citation: Cruddace et al. 2002
Citation function: Uses
Since the ground-breaking observation of radio pulsars by Hewish et al. (1968), NSs remain one of the most studied astrophysical objects; they are assumed to be possible sources of high-energy emission. The typical values of the surface magnetic field as inferred from simple magnetic dipole models and spin-down rates are in the range 108–1013 G (Taylor et al. 1993; Alpar et al. 1982). Note that among radio pulsars, PSR J1847−0130 exhibits a strong magnetic field of B = 9.4 × 1413 G (McLaughlin et al. 2003). However, besides the X-ray luminosities observed from the anomalous X-ray pulsars (AXPs), the inferred periods of AXPs and soft-γ repeaters suggest that such NSs have even larger surface magnetic fields of 1014−1015 G (Duncan & Thompson 1992; Paczynski 1992; Thompson & Duncan 1996; Melatos 1999). NSs with such high surface magnetic fields are popularly known as magnetars. Several interesting studies (Usov 1992; Kluźniak & Ruderman 1998; Wheeler et al. 2000; Starling et al. 2009; Cenko et al. 2010) have predicted that magnetars are the probable source of γ-ray bursts, and they require a higher magnetic field such as 1016–1017 G to initiate Poynting-flux-dominated jets. Although until now only around 30 magnetars have been detected, researchers have speculated that these astrophysical objects may account for 10% of the NS population (Kouveliotou et al. 1998). For NSs, the effects of a strong magnetic field on the ultra-dense electron gases in their interiors have been studied in several papers (Canuto & Ventura 1977; Fushiki et al. 1989; Abrahams & Shapiro 1991; Fushiki et al. 1992; Roegnvaldsson et al. 1993). Studies of dense and strongly magnetized nuclear matter have also been carried out by Chakrabarty et al. (1997), Bandyopadhyay et al. (1998), Broderick et al. (2000), Suh & Mathews (2001), Harding & Lai (2006), Chen et al. (2007), Rabhi et al. (2008), and references therein. Finally, we mention that studies of NSs with different magnetic-field configurations, viz., toroidal, poloidal, or mixed were carried out by Bocquet et al. (1995), Cardall et al. (2001), and Pili et al. (2014).
Citation: Paczynski 1992
Citation function: Background
As we discussed in the introduction, different choices of age anchors for the 26Al–26Mg decay system, CAIs versus the D’Orbigny angrite, result in different ages with a difference of 1–1.5 Ma. The differences can be interpreted as a reflection of different initial 26Al/27Al in the CAI accretion region, closest to the Sun (MacPherson, Wark & Armstrong 1988; Sossi et al. 2017) compared with the initial 26Al/27Al in the angrite formation region. In fact, CAIs also show large mass-independent isotopic anomalies for multiple elements, e.g. O, Cr, and Ti (Trinquier et al. 2009; Krot et al. 2020), which suggest that they formed in different nebula environments compared to achondrites like angrites and EC 002. In Fig. 3 and Table S6, we show U–Pb, 53Mn–53Cr, and 26Al–26Mg age comparisons for several achondrites. The data indicates that for all NC-like achondrites [ε54Cr values  0.3; Zhu et al. (2021a)], their U–Pb ages, and 53Mn–53Cr and 26Al–26Mg ages anchored to D’Orbigny are consistent. This consistency may indicate that the abundance of initial 26Al/27Al value of NC achondrites at 4567 Ma may have been lower in the non-carbonaceous region relative to the canonical initial Solar system value of 26Al/27Al derived from CAIs (Schiller et al. 2015). However, for NWA 6704, which has a CR chondrite-like ε54Cr, its 53Mn–53Cr ages are more consistent if its 26Al–26Mg age is anchored to CAIs, which suggests that the 26Al abundance of the CC bodies in the outer Solar system may follow that of CAIs. This conclusion is also consistent with the olivine grains (with Al/Mg ratio close to 0) in carbonaceous chondrites having the canonical initial μ26Mg (∼−35 ppm), indicating that the initial 26Al abundance in CCs at 4567 Ma was similar to that of CAIs (Gregory et al. 2020). Thus, the age difference of 26Al–26Mg ages relative to 53Mn–53Cr ages can be caused by a different initial abundance of 26Al in the inner relative to the outer Solar system. We note that the U–Pb age of NWA 2976 overlaps the Al–Mg ages anchored to both CAIs and D’Orbigny. This issue warrants more detailed assessment and discussion, e.g. the U–Pb dating (Krestianinov, Datta & Amelin 2021); however, mineral fractions used for dating should avoid xenocrystic material that may have formed at different times and from different sources.
Citation: Trinquier et al. 2009
Citation function: Background; Compare/Contrast
The new sample of HSPs gave us an opportunity to see the optical IDVs of HSPs and compare its properties with optical IDVs of LSPs. We started a dedicated project to search for optical IDV in HSPs and after doing 62 nights of IDV observations of HSPs which gave us 144 LCs (41 in B band, 62 in R band, and 41 in B−R colour) of five HSPs (Mrk 421, 1ES 1426+428, 1ES 1553+113, 1ES 1959+650, and 1ES 2344+514). Interestingly, we found that, four HSPs did not show any IDV (Gaur, Gupta & Wiita 2012a; Gaur et al. 2012b,c), but only one HSP 1ES 1426+428 for which we have the least observations have shown IDV in six LCs out of eight LCs (Gaur et al. 2012c). Our this pilot project gave us 6 IDV LCs out of 144 LCs searched for IDV i.e ∼4 per cent LCs have shown IDV. We explained it by density inhomogeneities and bends in the bases of the jets by Kelvin–Helmholtz instabilities (Romero, Cellone & Combi 1999). We gave an alternative explanation i.e. since in HSPs, the optical band lies below the SED peak, hence, we should see changes in the efficiency of acceleration of, and/or in the rates at which energy is radiated by, the highest energy electrons available for synchrotron emission would have a more retarded effect on optical variability in HSPs (Gaur et al. 2012b). In LSPs, the optical band is dominated by highest energy electrons emitting synchrotron radiation and probably the X-ray emission is dominated by the comparatively lower energy electrons emitting the inverse Compton radiation, hence their X-ray variability is less pronounced than optical variability. If SED peak is really responsible for IDV properties, then we suspected that X-ray IDV LCs in LSPs should not show any IDV at all or show on rare occasions. With this motivation, here we present the X-ray IDV study of almost complete sample of 10 LSPs and 2 ISPs observed by XMM–Newton since its launch and we found that the LSPs show very less IDV 2 out of 50 LC i.e. 4 per cent in X-ray bands. We have reported above the similar finding for HSPs in optical bands.
Citation: Gaur et al. 2012c
Citation function: Compare/Contrast
Many classical theories are presented exploring the dynamics of various rotating plasma instabilities, such as those of Vedenov et al. (1961) and Rosenbluth & Simon (1965). Initially studied by Velikov (1959), work in this area starts with the nondissipative Couette flow of an ideal conducting fluid between two differentially rotating cylinders. It is observed that the fluid is destabilized by applying an axial magnetic field. This work was confirmed by Chandrasekhar (1960). By now, research on magnetorotational instability (MRI) is well understood and reported in both experimentally (liquid metals; Rüdiger et al. 2003; Ji et al. 2004) and theoretically based studies of astrophysical plasmas (Blaes & Balbus 1994; Hawley et al. 1995; Quataert et al. 2002). In astrophysical environments, MRI is now believed to be the main factor deriving plasma turbulence and also to account for the accretion process that feeds stars and black holes (Kotake et al. 2004). After its discovery, MRI remained outside the general theory of plasma instabilities for almost four decades, and was revisited with its application to astrophysical objects by considering the anomalous viscosity problem in the accretion disks surrounding astrophysical compact objects; see Balbus & Hawley (1991). MRI is supposed to be a destructive factor for the hydrodynamical Rayleigh’s stability criterion, making it magnetohydrodynamically unstable and leading directly to orbital angular momentum transport and disk turbulence (Hawley & Balbus 1991, 1992; Balbus & Hawley 1998). A two-fluid model for MRI in differentially rotating plasma was presented by Ren et al. (2011) to describe the dynamical behavior of ions and electrons in both low- and high-frequency regimes. They have demonstrated the instability criteria that differ significantly from the one derived in the single-fluid model (Balbus & Hawley 1991). Mikhailovskii et al. (2008a) studied MRI in dusty plasmas using a single-fluid MHD model. They derived a generalized relation by including the effect of dust particulates, and due to the induced electric field the rotation of equilibrium plasma was modified and gave rise to an instability called a dust-induced rotational instability. In their preceding work (Mikhailovskii et al. 2008b) with an immobile dust grain assumption, a comprehensive analysis was carried out for the collective phenomenon of instability which provides a more detailed theory about MRI. The MRI in differential rotating dusty plasmas with dissipative effects was investigated by using local linear analysis with effect of immobile dust grains (Ren et al. 2009).
Citation: Kotake et al. 2004
Citation function: Background
During the X-ray ‘hard’ state, the radio and X-ray behaviour of BHXBs is correlated and has been studied in depth for many sources using quasi-simultaneous observations (e.g. Corbel et al. 2003; Gallo, Fender & Pooley 2003; Coriat et al. 2011; Corbel et al. 2013), and is known as the radio:X-ray plane. It was thought that all XBs followed a relation in the form of $L_{\rm Radio} \propto L_{\rm X-ray}^{0.6}$, based upon early observations of GX 339−4 (Hannikainen et al. 1998; Gallo et al. 2003; Corbel et al. 2003, 2013). This relation extends down to very low luminosities, i.e. into quiescence (Corbel et al. 2003, 2013; Plotkin et al. 2017; Tremou et al. 2020), and has been observed in other sources such as V404 Cygni (e.g. Corbel, Koerding & Kaaret 2008) and XTE J1118+480 (Gallo et al. 2014). However, further observations revealed the presence of another population of BHXBs which are less radio luminous than this relation, a so-called ‘radio-quiet’ branch, which followed $L_{\rm R} \propto L_{\rm X}^{1.4}$, such as H 1743−322 (e.g. Coriat et al. 2011; Williams et al. 2020). In some of these ‘radio-quiet’ objects, they are then found to re-join the ‘radio-loud’ branch when they go back into quiescence (Coriat et al. 2011; Carotenuto et al. 2021). The underlying cause of the split tracks (see Gallo, Miller & Fender 2012; Gallo et al. 2014; Gallo, Degenaar & van den Eijnden 2018, for a clustering analysis into the statistical robustness of this split) for BHXBs is not known, but it may be due to differences in the radiative efficiency of the accretion flow (Coriat et al. 2011; Koljonen & Russell 2019), an inclination effect of the source (Motta, Casella & Fender 2018), differences in the accretion disc contribution (Meyer-Hofmeister & Meyer 2014) or changes in the magnetic field (Casella & Pe’er 2009). For the purposes of this paper, we will refer to the original $L_{\rm R} \propto L_{\rm X}^{0.6}$ correlation sources as ‘radio-loud’ objects, and those that diverge on to the $L_{\rm R} \propto L_{\rm X}^{1.4}$ track as ‘radio-quiet’ sources.
Citation: Gallo et al. 2014
Citation function: Background; Background; Background