adsabs/FOCAL huggingface dataset:
The column S in the table shows that at least two minima of line-of-sight velocity irregularity are present for all numbers m considered: the first short-wavelength minimum with the radial scale λrad ≈ 1 kpc and the second long-wavelength minimum with the scale λrad ≈ 2.5 kpc. There are three possible ways to explain these scales. In the first place, one can identify the scales as fully independent global modes. This is because for a fixed number of spiral arms m, in a self-gravitating rotating astrophysical disc configuration, there usually exists a sequence of radial wavelengths (Bertin et al. 1977; Griv & Chiueh 1998; Griv et al. 2000, 2006, Griv et al. 2012; Fujii et al. 2011; Griv & Wang 2014). The non-negative azimuthal mode number m gives the number of spiral arms, whereas the real radial wavenumber krad = 2π/λrad determines the radial structure.4 Secondly, one can identify these scales as Fourier harmonics of a single wave (Mishurov et al. 1979, p. 151). And thirdly, the phase ϕ0 that we found is between −48° and −70° and/or between 327° and 351° for all m. To be consistent with the current ideas as to the Galactic spiral structure (Vallée 2008, fig. 1 therein; Xu et al. 2013, fig. 12 therein; Hou & Han 2014, fig. 5 therein; Nakanishi & Sofue 2016, fig. 7 therein), we may interpret this finding to mean that the kinematics of 223 Cepheids investigated reveals the segments of different large-scale arms. Namely, whatever spiral m models we consider, the Sun is located at the outer edge of a local spiral arm (at the outer edge of the Cygnus–Orion arm; ϕ0 ≈ −60° for the first minimum) and/or at the outer edge of an inner spiral arm (at the outer edge of the Carina–Sagittarius arm, very close to the arm centre; ϕ0 ≈ 330° for the second minimum) (for notation, see Reid et al. 2009, fig. 1 therein; Choi et al. 2014, fig. 14 therein; Sanna et al. 2014, fig. 6 therein; Camargo et al. 2015, fig. 12 therein; Sato et al. 2014, fig. 3 therein). Basharina et al. (1980) have already confirmed the Sun to be located in the local arm. Further studies are needed to clarify the issue (see also Hou & Han 2014).
Citation: Griv & Wang 2014
Citation function: Uses
Disk galaxies typically reveal out-of-equilibrium features due to incomplete equilibration. These may appear in the form of bars and spiral arms, which are large-scale perturbations in the radial and azimuthal directions, responsible for a slow, secular evolution of the disk. In the vertical direction, disks often reveal warps (Binney 1992). In the case of the Milky Way (MW) disk, which can be studied in much greater detail than any other system, recent data from astrometric and radial velocity surveys such as SEGUE (Yanny et al. 2009), RAVE (Steinmetz et al. 2006), GALAH (Bland-Hawthorn et al. 2019), LAMOST (Cui et al. 2012), and above all Gaia (Gaia Collaboration et al. 2016, 2018a, 2018b) have revealed a variety of additional vertical distortions. At large galactocentric radii (>10 kpc) this includes, among others, oscillations and corrugations (Xu et al. 2015; Schönrich & Dehnen 2018), and streams of stars kicked up from the disk that undergo phase mixing, sometimes referred to as “feathers” (e.g., Price-Whelan et al. 2015; Thomas et al. 2019; Laporte et al. 2022). Similar oscillations and vertical asymmetries have also been reported in the solar vicinity (e.g., Widrow et al. 2012; Williams et al. 2013; Yanny & Gardner 2013; Gaia Collaboration et al. 2018b; Quillen et al. 2018; Bennett & Bovy 2019; Carrillo et al. 2019). One of the most intriguing structures is the phase-space spiral discovered by Antoja et al. (2018) and studied in more detail in subsequent studies (e.g., Bland-Hawthorn et al. 2019; Li 2021; Li & Widrow 2021; Gandhi et al. 2022). Using data from Gaia DR2 (Gaia Collaboration et al. 2018a), Antoja et al. (2018) selected ∼900,000 stars within a narrow range of galactocentric radius and azimuthal angle centered around the Sun. When plotting the density of stars in the (z, v

z
)-plane of vertical position, z, and vertical velocity, v

z
, they noticed a faint, unexpected spiral pattern, which became more enhanced when color-coding the (z, v

z
)-“pixels” by the median radial or azimuthal velocities. The one-armed spiral makes two to three complete wraps, resembling a snail shell, and is interpreted as a signature of phase mixing in the vertical direction following a perturbation, which Antoja et al. (2018) estimate to have occurred between 300 and 900 Myr ago. More careful analyses in later studies (e.g., Bland-Hawthorn et al. 2019; Li 2021) have nailed down the interaction time to ∼500 Myr ago.
Citation: Cui et al. 2012
Citation function: Motivation; Motivation; Motivation
Brown dwarfs and giant exoplanets populate the same temperature range and share many physical properties, such as their molecule-dominated atmospheres and gradual cooling from ~3000 K at formation to ~100 K like the solar system gas-giant planets. Recent discoveries of very massive planets (Chauvin et al. 2005; Marois et al. 2010; Delorme et al. 2013), some possibly more massive than the 13 MJup deuterium burning mass limit, hint that planets could overlap with brown dwarfs in mass. On the other hand, the discovery of isolated L dwarfs in young clusters (Zapatero Osorio et al. 2002, 2014; Peña Ramírez et al. 2012), in young moving groups (Liu et al. 2013; Gagné et al. 2015; Gauza et al. 2015), and very cold very nearby Y dwarf objects (e.g., Kirkpatrick et al. 2012; Luhman 2014) show that very low-mass isolated brown dwarfs exist and overlap with the planetary masses. When these low-mass brown dwarfs are close enough and bright enough to be observed spectroscopically their atmospheres are much easier to study than similar exoplanets that lie near their very bright host stars. Liu et al. (2013) notably showed that the ~8 MJup brown dwarf PSO J318.5−22, a β-pictoris moving group member shares the spectral characteristics of the young directly imaged exoplanets, as well as atypically red late-L spectral type objects (e.g., Faherty et al. 2013; Gizis et al. 2015; Kellogg et al. 2016; Schneider et al. 2014, 2016; Bonnefoy et al. 2016). When CFBDSIR J214947.2−040308.9, hereafter CFBDSIR 2149, was identified (Delorme et al. 2012), it seemed to be a candidate member of the AB Doradus young moving group and, together with the low-gravity features in its spectrum, made it a unique T-type isolated planetary-mass candidate. Another earlier-type, isolated young planetary-mass T-dwarf, SDSS J111010.01+011613.1, has been identified as a bona fide member of AB Doradus moving group (\hbox{$149^{+51}_{-19}$}149-19+51 Myr; Bell et al. 2015) by Gagné et al. (2015). The late-T spectral type of CFBDSIR 2149 is typical of the coolest known directly imaged exoplanets, such as GJ 504 b or 51 Eri b (Kuzuhara et al. 2013; Macintosh et al. 2015), that the latest generation of adaptive optics systems are detecting. We therefore carried out a multi-wavelength, multi-instrument follow-up of CFBDSIR 2149 to fully characterise it and constrain its nature.
Citation: Ramírez et al. 2012
Citation function: Motivation; Motivation
FRBs have been detected in the radio band. Circular polarization of FRB 140514 was measured at a level of 





 degree, while linear polarization is less than 10% as a 1σ upper limit (Petroff et al. 2015). We apply our model to self-consistently get both circular and linear polarization degrees of FRB 140514. The case of B = 1010 G, Te = 108 K, and γth = 104 shown as the solid line in the right panels of Figure 3 seems to be one solution to explain the polarization observations. Thus, we may consider FRB 140514 to be associated with strongly magnetized sources. A neutron star can be one candidate for FRB origin. Here, we assume that the polarized emission is pointing toward observers, as usual neutron star models assumed, but we do not include additional jet view-angle effects. We note one recent theoretical work on the polarization detection of neutron stars in the soft X-ray band (Yatabe & Yamada 2017), and the relatively low electron temperature from our analysis indicates that the radiation of FRB 140514 may not occur near the neutron star magneto-atmosphere above the neutron star surface. On the other hand, the alternative case of B = 1 G, Te = 1010 K, and 





 in our model can obtain the circular polarization degree of 20%. However, in this case, the intrinsic linear polarization degree is about 76%. This seems to be in contrast with the observation of FRB 140514. Therefore, we are inclined to believe from our polarization analysis that FRB may have a neutron star origin. Besides FRB 140514, FRB 110523 has the linear polarization degree of 





, but the circular polarization degree of about 23% due to instrumental leakage was not confirmed to be an astronomical result (Masui et al. 2015). FRB 081507 has the high linear polarization degree of 80% (Ravi et al. 2016) and FRB 150215 has the linear polarization degree of 





 (Petroff et al. 2017). If we can detect both linear and circular polarization for FRBs, we may further constrain the FRB physics.
Citation: Yatabe & Yamada 2017
Citation function: Compare/Contrast; Compare/Contrast
To this we can add that Zimbardo and coworkers reported in a series of papers that spacecraft observations of SEP intensity-time profiles ahead of interplanetary shocks yielded power laws instead of the exponential profiles predicted by standard DSA. They interpreted this as a sign of superdiffusive spatial SEP transport in the shock vicinity resulting in superdiffusive shock acceleration. They predicted the superdiffusive shock acceleration to be more efficient, producing a harder accelerated power-law spectrum for SEPs, than standard DSA (Zimbardo et al. 2006; Perri & Zimbardo 2007, 2008, 2009; Zimbardo & Perri 2013; Zimbardo et al. 2017, 2020; Zimbardo & Perri 2020). Zimbardo & Perri (2020) explained such superdiffusive transport as a product of subdiffusion in pitch-angle space arising from the observation that Alfvén-wave pitch-angle scattering times in the solar wind have a power-law distribution instead of a single value. Another possibility is that there is a strong presence of SMFRs in the vicinity of heliospheric shocks so that particle trapping effects could result in anomalous diffusive transport and shock acceleration. Evidence is increasing that, in a space plasma with a significant guide/background magnetic field like the solar wind, SMFRs naturally form as part of a self-generated, quasi-two-dimensional (quasi-2D) MHD turbulence component that might dominate other MHD wave turbulence modes. This is supported by observations in the slow solar wind near 1 au (e.g., Matthaeus et al. 1990; Bieber et al. 1996; Greco et al. 2009; Zheng & Hu 2018), MHD simulations (e.g., Shebalin et al. 1983; Dmitruk et al. 2004), and nearly incompressible MHD turbulence transport theory (e.g., Zank & Matthaeus 1992, 1993; Zank et al. 2017, 2018, 2020). It is plausible that SMFRs can have a strong presence ahead of interplanetary shocks. Furthermore, there is evidence that the occurrence of SMFRs peaks at large-scale current sheets in the solar wind where magnetic reconnection produces additional SMFRs that can trap energetic particles (Khabarova et al. 2015, 2016; Hu et al. 2018). Many of these current sheets occur behind traveling shocks (Khabarova & Zank 2017). Recently, a theoretical investigation was launched to determine the transmission of quasi-2D SMFRs through a perpendicular shock (Zank et al. 2021). The results show a strong enhancement in the magnetic energy density of the SMFR magnetic island component behind the shock, thus confirming a strong presence of SMFRs downstream of the shock as well.
Citation: Zimbardo & Perri 2013
Citation function: Compare/Contrast
The first SMGs were detected using SCUBA at 850 μm (Smail et al. 1997; Barger et al. 1998; Hughes et al. 1998), which remains one of the prime wavelengths to detect these galaxies (e.g., Geach et al. 2017), thanks to a combination of available instruments, spectral window, and the negative k-correction at that wavelength. Other single-dish samples of SMGs have also been obtained at 1.1–1.3 mm using MAMBO (e.g., Eales et al. 2003; Bertoldi et al. 2007; Greve et al. 2008) and AzTEC (e.g., Aretxaga et al. 2011; Yun et al. 2012), at 1.4 mm/2 mm with the SPT (Vieira et al. 2010), and at 2 mm with GISMO (Staguhn et al. 2014; Magnelli et al. 2019). Selecting SMGs from observations at longer wavelengths is thought to favor galaxies at higher redshifts (e.g., Smolčić et al. 2012; Vieira et al. 2013; Staguhn et al. 2014; Magnelli et al. 2019; Hodge & da Cunha 2020), although it is difficult to compare the redshift distributions in an unbiased way (see, e.g., Zavala et al. 2014 for a discussion), and account for intrinsic variations of galaxy far-IR spectral energy distributions (SEDs). Nevertheless, the 2 mm band has been put forth as a potential candidate to detect high-redshift (z > 3) galaxies (e.g., Casey et al. 2018a, 2018b, 2019; Zavala et al. 2021). The negative k-correction is stronger at 2 mm than at 850 μm; thus, for a fixed SED, the 2 mm band should pick up more high-redshift galaxies than at 870 μm. In addition, better atmospheric transmission and larger fields of view can be achieved at 2 mm (but corresponding poorer resolution). Such an effort is currently ongoing (see Zavala et al. 2021 for first results). To understand the relationship between the populations detected at 850 μm and at 2 mm, we require a detailed characterization of the (sub)millimeter SEDs of these sources. Multiwavelength submillimeter observations are still rare, with most observations focusing on a single wavelength. Only a handful of sources observed at 2 mm have complementary shorter-wavelength detections (Staguhn et al. 2014; Magnelli et al. 2019). Thus, a more systematic multiwavelength dust continuum investigation is warranted in order to reveal the dust properties of (sub)millimeter-detected sources.
Citation: Staguhn et al. 2014
Citation function: Motivation; Motivation
To test IADI, we made a model of a disk in scattered light. This model is 400 × 400 pixels, each pixel having a physical size of 1 × 1 au. Multiple rings are implemented, inspired by observations of disks (e.g., HD 97048 Ginski et al. 2016; RX J1615 de Boer et al. 2016 and TW Hydrae van Boekel et al. 2017), see panel a in Fig. 2. A specific inclination and rotation is achieved using the warpAffine function from OpenCV (Bradski 2000). Young planet-forming disks are still gas rich and dust particles are stratified due to gas pressure along the vertical axis. Therefore, flaring is implemented in the model via an offset of the rings with respect to the center of the ring (see the detailed discussion in de Boer et al. 2016). For our disk model, we are using the power-law profile for the scattering surface height H and the separation r found by Ginski et al. (2016) for the disk around HD 97048. This profile describes the flaring discovered in this source reasonably well up to a separation of ~270 au. Considering that the model disk used will have a separation of 200 au, this formula will be sufficient to simulate the disk height2. Illumination effects of the central star are implemented via a ~1/r2 intensity dependence from the center of the disk, making the inner part of the disk brighter compared to the outer part. Moreover, the intensity also depends on the light scattering angle via the phase function. Because a physical model is beyond the scope of this work, a “pseudo” phase function is implemented to mimic the same asymmetries in light distribution seen in observed disks via I = cos ϕ, where the intensity I depends on the cosine of the azimuthal angle ϕ. This pseudo phase function depends on the azimuthal angle instead of a scattering angle on which a real phase function would depend. This makes the part of the disk facing toward the observer appear brighter than the part facing away in a fairly simple way. Lastly, the model is put through a Gaussian convolution kernel from the scipy ndimage package (Virtanen et al. 2020) to remove sharp edges and give a finite resolution to the model. The final three steps are shown in the top row of Fig. 2.
Citation: Ginski et al. 2016
Citation function: Uses; Uses
In Section 3.2, it was noted that the long range structure observed in the reaction between the formamide and amino radicals was that of neutral hydrogenated isocyanic acid. However, instead of hydrogenation, protonation could be occurring. Isocyanic acid (HNCO) has been detected in the ISM and can be formed from simple building blocks like atomic nitrogen and CO (Snyder & Buhl 1972; Buhl, Snyder & Edrich 1973; Nourry, Zins & Krim 2015). It has been used in the study of the formation of urea in irradiation studies, where its irradiation is shown to lead to the formation of formamide and urea (Ferus et al. 2018; Raunier et al. 2003, 2004). In addition, it has been shown to be a product of the dissociation of formamide and has even been suggested to be a catalyst for the formation of H2 in the ISM (Lundell, Krajewska & Räsänen 1998; Duvernay et al. 2005; Haupa, Tarczay & Lee 2019). Moreover, there is a clear correlation with the abundances of formamide (Bisschop et al. 2007; López-Sepulcre et al. 2015). There is ample evidence of HNCO molecules functioning as an imine base, both in the gas phase and in the solid state (Dekock & Jasperse 1983; Hop et al. 1989; Hunter & Lias 1997; Hudson, Khanna & Moore 2005; Gupta et al. 2013; Bouchoux 2018; Marcelino et al. 2018). The proton affinity of HNCO is higher than that of water, so that one would expect protonation to occur even in water-based ices (Dekock & Jasperse 1983; Bouchoux 2018). Thus, in this section, calculations starting from protonated isocyanic acid are reported. Only DFT calculations were performed on these reactions, since multiconfigurational effects can be assumed to be small for these non-radical charged species. The first step that was investigated is the protonation of isocyanic acid. There are multiple options for protonation in this case: Protonation can occur on the nitrogen, carbon, or oxygen atoms. Hereby, we note that N-protonated isocyanic acid has been detected in laboratory experiments and in the ISM in the dense core L483 with a more tentative detection towards Sgr B2(N) (Hudson et al. 2005; Gupta et al. 2013; Marcelino et al. 2018). Moreover, O-protonated isocyanic acid has also been detected in laboratory experiments (Lattanzi et al. 2012). The potential energy curves for protonating on each of these atoms is shown in Fig. 7. Protonation on the nitrogen and oxygen atoms is barrierless. The well depths are 603.1 and 603.9 kJ mol−1, respectively, making them equally likely. Protonation on the carbon atom, however, is not barrierless, but requires an energy input of 14.4 kJ mol−1. It also results in a much less stable species, with a well depth of only 172.8 kJ mol−1, making it the least likely of the three pathways. Nitrogen protonation is the product that most naturally leads on to the formation of urea as it creates one of the amino groups found in urea. None the less, the similarity in energies for the nitrogen- and oxygen-protonated species means that the oxygen-protonated species could also be formed. The latter is the species we will consider first.
Citation: Bisschop et al. 2007
Citation function: Background
In a quasi-isothermal jet, Uj is (17)\begin{equation} \label{eq:U_j_quasi} U_{\rm j} = \zeta n_0 m_{\rm p} c^2\left(\frac{\gamma_{\rm j}\beta_{\rm j}}{\gamma_0\beta_0}\right)^{-\Gamma}\left(\frac{z}{z_0}\right)^{-2} \cdot \end{equation}Uj=ζn0mpc2γjβjγ0β0−Γzz0-2·Substituting Eqs. (17) and (13) into Eq. (10), and assuming the jet is launched with an initial γ0β0 equal to the sound speed (Eq. (16)), the 1D Euler equation that results is \begin{eqnarray} \label{eq:AGNJET_Corrected} &amp;&amp;\left\{\gamma_{\rm j}\beta_{\rm j}\frac{\Gamma+\xi}{\Gamma-1}-\Gamma\gamma_{\rm j}\beta_{\rm j}-\frac{\Gamma}{\gamma_{\rm j}\beta_{\rm j}}\right\}\frac{\partial \gamma_{\rm j}\beta_{\rm j}}{\partial z} = \frac{2}{z}; \\ &amp;&amp;\xi = \frac{1}{\zeta}\left(\frac{\gamma_{\rm j}\beta_{\rm j}}{\gamma_0\beta_0}\right)^{\Gamma-1}; \qquad \gamma_0\beta_0=\sqrt{\frac{\zeta\Gamma(\Gamma-1)}{1+2\zeta\Gamma-\zeta\Gamma^2}} \cdot \end{eqnarray}γjβjΓ+ξΓ−1−Γγjβj−Γγjβj∂γjβj∂z=2z;ξ=1ζγjβjγ0β0Γ−1; γ0β0=ζΓ(Γ−1)1+2ζΓ−ζΓ2·The above equation should reduce to the jet Lorentz factor profile used in Falcke (1996), Markoff et al. (2005) when ζ = 1. However, it differs from Eq. (2) in Falcke (1996): (20)\begin{equation} \label{eq:Heino96} \left\{\gamma_{\rm j}\beta_{\rm j}\frac{\Gamma+\xi}{\Gamma-1}-\frac{\Gamma}{\gamma_{\rm j}\beta_{\rm j}}\right\}\frac{\partial \gamma_{\rm j}\beta_{\rm j}}{\partial z} = \frac{2}{z}; \end{equation}γjβjΓ+ξΓ−1−Γγjβj∂γjβj∂z=2z;(21)\begin{equation} \xi = \left(\gamma_{\rm j}\beta_{\rm j}\frac{\Gamma+1}{\Gamma(\Gamma-1)}\right)^{1-\Gamma} \cdot \end{equation}ξ=γjβjΓ+1Γ(Γ−1)1−Γ·The difference between our equation and the equation in Falcke (1996) can be accounted for as follows: the − Γγjβj term in Eq. (18) results from a neglected \hbox{$\frac{\partial}{\partial z}(U_{\rm j}/n)$}∂∂z(Uj/n) term, the difference in the exponent in ξ results from an arithmetic error, and finally the difference in the inside of the parenthesis of ξ terms is from setting \hbox{$\gamma_0\beta_0 = \beta_{\rm s0}^{2}$}γ0β0=βs02 instead of using the proper value given in Eq. (16). The difference between the solutions of Eqs. (18) and (20) are small and shown in Fig. 1. In Fig. 1, we also include solutions to the 1D Euler equations when the jet is isothermal (Tj = const., i.e., Eq. 20 with ξ = 1) and adiabatic (Tj ∝ (γjβj)1 − Γz2 − 2Γ, see Eq. (25)).
Citation: Falcke (1996)
Citation function: Compare/Contrast; Compare/Contrast