adsabs/FOCAL huggingface dataset:
Planet formation occurs in disks around young stellar objects. Interactions between planets and disks are very complex. Young planets are expected to cause rings, cavities, spirals, and disturbances in the velocity field and other features in the disk, which in turn may be used to infer the presence of these young planets. In the past few years, much evidence about this phase of planet formation has been accumulated because high-resolution images in the millimeter and sub-millimeter wavelength ranges have been provided by the Very Large Array (VLA) and the Atacama Large Millimeter Array (ALMA; see e.g. the case of HL Tau; ALMA Partnership 2015), and by high-contrast imagers such as the Gemini Planet Imager (GPI, Macintosh et al. 2014) and SPHERE (Spectro- Polarimetic High contrast imager for Exoplanets REsearch, Beuzit et al. 2008; see, e.g., Avenhaus et al. 2018). The literature on indirect evidence of the presence of planets is now becoming very rich, and nearby young stars surrounded by gas-rich disks are intensively studied for this purpose. In most cases, available data cannot fully eliminate alternative hypotheses, or the data have ambiguous interpretations (see, e.g., Bae et al. 2018 and Dong et al. 2018), although strong indirect evidence of the presence of planets from local disturbances of the velocity field have recently been considered for the case of HD 163296 (Pinte et al. 2018; Teague et al. 2018). In general, small grains are thought to be more strongly coupled with gas and are thus less sensitive to radial drift and concentration that can strongly affect large grains (see the discussion in Dipierro et al. 2018). For this reason, observations at short wavelengths provide an important complementary view of what can be seen with ALMA. On the other hand, a direct detection of still-forming planets embedded within primordial gas-rich disks, which is expected to be possible with high-contrast imaging in the near infrared (NIR), is still scarce; remarkable cases are LkCa-15 (Kraus & Ireland 2012; Sallum et al. 2015) and PDS-70 (Keppler et al. 2018; Müller et al. 2018; Wagner et al. 2018). In particular, in this second case, a clear detection of an accreting planet in the cavity between the inner and outer ring was obtained, making it an archetype for planet formation and planet-disk interactions. However, many cases remain ambiguous; a classical example is HD 100546 (see, e.g., Quanz et al. 2013a, 2015; Currie et al. 2014, 2015; Rameau et al. 2017; Sissa et al. 2018).
Citation: Dong et al. 2018
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Figure 10 (see also Figs. A.4 and A.5) compare the spatial distributions of the Hα and 857 GHz emission in the Taurus–California–Perseus region (e.g., Taurus, Auriga, California, and Perseus). The 857 GHz dust emission traces each molecular cloud and exhibits a hole-like structure. This hole-like structure can also be seen in HI emission, as shown in Figs. A.4 and A.5. The Hα emission fills the hole-like structure seen in the 857 GHz dust emission near the center of the field. The Taurus, California, and Perseus molecular complexes traced by the 857 GHz dust emission are distributed at the edge of the hole-likestructure. Lim et al. (2013) also found evidence of a shell-like structure using dust extinction and 12 CO (1–0) maps. The hole-like structure may result from the expansion of a large-scale supershell produced by a supernova in the Per OB2 association that compresses the Taurus cloud from the far side (Olano & Poeppel 1987; Bally et al. 2008). An Hα absorption feature is detected toward the Taurus cloud (see Figs. 10 and A.6), suggesting that the Taurus cloud lies at the front surface of the large-scale supershell produced by the Per OB2 association. The distance to the Per OB2 association is estimated to be 340 pc from the Sun (Cernis 1993), while the distance to the Taurus cloud is ~140 pc (Elias 1978). These distances are consistent with the Taurus cloud lying in front of the Per OB2 association. The B211/B213 filament also appears to be in front of the HI shell (see Fig. 10 in Chapman et al. 2011). This morphology suggests that the B211/B213 filament may have formed as a result of an expanding supershell. This may provide another reason for the different initial gas velocities for the northeastern and southwestern sheet components in addition to large-scale acceleration by the gravitational potential of the B211/B213 cloud (see Sect. 4.2.3). The Local Bubble surrounding the Sun might also compress the Taurus cloud from the opposite direction. The Local Bubble surrounding the Sun was produced by supernovae (Snowden et al. 1998; Sfeir et al. 1999), and the wall of the Local Bubble is located close to the Taurus cloud (Könyves et al. 2007; Lallement et al. 2014).
Citation: Olano & Poeppel 1987
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Assuming that the force-field approach to the solution of the Parker (1965) cosmic-ray transport equation is valid, the connection between historic cosmic-ray intensities and the solar properties they encountered lies in the effective diffusion coefficient that is assumed in this approximation. Establishing such a connection, however, is no simple task. Many theories have been proposed to describe the scattering of cosmic rays in the heliosphere. The most likely candidates for this task, given their reasonable agreement with observations and numerical simulations of cosmic-ray diffusion coefficients, are the quasilinear theory (QLT) of Jokipii (1966), the weakly nonlinear theory (WNLT) of Shalchi et al. (2004b), and the nonlinear guiding center theory (NLGC; or one of its variants; see, e.g., Matthaeus et al. 2003; Shalchi 2006, 2009, 2010; Ruffolo et al. 2012). Shalchi (2009) provides in depth theoretical treatments of most the abovementioned theories. These scattering theories all require as a key input an expression for the power spectrum of the turbulent fluctuations of the HMF. These spectra depend upon basic turbulence quantities, such as the magnetic variance, and various correlation scales. Turbulence power spectra are discussed in detail by, e.g., Batchelor (1970) and Matthaeus et al. (2007), whereas more background on the abovementioned turbulence quantities can be found in, e.g., Matthaeus & Goldstein (1982), Petrosyan et al. (2010), Matthaeus & Velli (2011), and Bruno & Carbone (2013). These basic turbulence quantities have been observed to show a marked dependence on the solar cycle at Earth (see, e.g., Smith et al. 2006b; Burger et al. 2014; Zhao et al. 2018). It follows then that mean free paths derived from these scattering theories would be expected to depend on the solar cycle as well, and several studies have reported such a dependence. Chen & Bieber (1993) find from an analysis of cosmic-ray anisotropies and gradients as observed by means of NMs, that larger mean free paths are associated with solar minima, and smaller mean free paths with solar maxima. The authors also report a mean free path dependence on solar magnetic polarity. Nel (2016) and Zhao et al. (2018) both extensively analyze spacecraft observations, using the turbulence quantities so calculated as inputs for expressions for diffusion coefficients derived from the QLT and NLGC theories. Both authors report that the resulting mean free paths display solar cycle dependences.
Citation: Matthaeus et al.; 2010
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Citation: Ruffolo et al. 2012
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The first investigations of the rotational spectra of ethyl isocyanide were carried out in 1966 by Bolton et al. (1966). The spectra of the first vibrational and torsional excited states were measured in the centimeter wave domain (Anderson & Gwinn 1968). In this initial study, the dipole moment was determined to be μa = 3.79 D and μb = 1.31 D; this value is usually large for a molecule that includes a CN group. This causes dense and intense rotational spectra in the millimeter wave range and also in the submillimeter wave range up to 900 GHz (bQ lines). Anderson & Gwinn (1968) also observed some A–E splittings due to the internal rotation motion of the methyl group. The most recent spectroscopic study is from Krüger & Dreizler (1992) who reinvestigated the internal rotation measurements and also determined hyperfine coupling parameters due to the nitrogen quadrupole. As in our previous studies of ethyl cyanide isotopologs, it was not possible to observe internal rotation and hyperfine splittings due to our Doppler limited resolution. Our analysis was rather easy, starting from a prediction based on Krüger & Dreizler (1992) parameters. First, we analyzed and fit the most intense transitions, the aRh transitions, up to 330 GHz. These transitions were shifted only a few MHz from the initial predictions. Then bR and bQ lines were searched and included in the fit up to 330 GHz. Next, all the spectra were analyzed up to 990 GHz without difficulty. For the fitting, we employed ASFIT (Kisiel 2001) and predictions were made with SPCAT (Pickett 1991). The global fits included 6 transitions from Anderson & Gwinn (1968), 29 lines from Krüger & Dreizler (1992), and 2906 from this work. The maximal quantum numbers are J = 103 and Ka = 30. Both reductions A and S were tested. A reduction permits us to check theagreement of our new parameters set with those from Krüger & Dreizler (1992) (Table 1). Using S reduction slightly decreases root mean square from 30.3 to28.7 kHz. The condition numbers are nearly the same: 295 and 310 for the A and S reductions, respectively.The A reduction requires 23 parameters, but 5 additional parameters are required for the S reduction (Table 2). For this reason we used the A reduction even if this molecule is close to the prolate limit with kappa = −0.9521. Part of the new measurements are in Table 3. Owing to its large size, the complete version of the global fit Table S1 is supplied at the CDS. The fitting files .lin (S2), .par (S3), and the prediction .cat (S4) are also available at CDS.
Citation: Krüger & Dreizler (1992)
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Citation: Krüger & Dreizler (1992)
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Citation: Krüger & Dreizler (1992)
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Regarding non-restricted hierarchical three-body systems, Krymolowski & Mazeh (1999) and Ford, Kozinsky & Rasio (2000) presented secular equations of motion (or Hamiltonian) up to the octupole order in semimajor axial ratio by using Hamiltonian perturbation techniques. Lee & Peale (2003) adopted both the octupole-level perturbation theory and direct numerical integrations to investigate the dynamical evolution for coplanar hierarchical planetary systems. In particular, the dynamics of apsidal resonance with critical argument of σ = ϖ1 − ϖ2 (ϖ1, 2 are the longitudes of pericentre) is studied and applied to some representative exoplanetary systems (Lee & Peale 2003). In a hierarchical planetary system with two comparable-mass planets orbiting a central star, Naoz et al. (2011) showed that orbits of the inner planet could flip from prograde to retrograde and back again due to the secular planet–planet interaction. Based on this behaviour, it becomes possible to form hot Jupiters on retrograde orbits by combining the eccentric ZLK effect and tidal friction (Naoz et al. 2011; Naoz, Farr & Rasio 2012; Teyssandier et al. 2013; Petrovich 2015; Petrovich & Tremaine 2016; Dawson & Johnson 2018). Naoz et al. (2013) re-derived the secular evolution equations for hierarchical three-body systems at the octupole-level approximation and found that orbital flips of inner planet are possible even at the quadrupole-level approximation. They pointed out that the relation h1 − h2 = π can be used to simplify the expression of Hamiltonian but the evolutions of H1 and H2 should be derived from the conservation of the total angular momentum rather than from the Hamiltonian canonical relations. Tan et al. (2020) explored the secular resonances with critical arguments arising in the Hamiltonian under the resonant Hamiltonian model, which is obtained by directly removing those terms involving short-period angles from the octupole-level Hamiltonian (i.e. only the secular and resonant terms are retained). It is of no problem when dealing with the quadrupole-order resonance (the ZLK resonance) because in this case the omitting terms are of octupole order. However, it may be inadequate to formulate the resonant model by directly removing those quadrupole-order periodic terms from the Hamiltonian when studying the octupole-order resonances. Hamers (2021) performed a semianalytic study about the properties of the ZLK oscillations at the quadrupole-level approximation, including the maximum eccentricities, time-scales of eccentricity/inclination oscillation and orbit flips. Naoz (2016) and Shevchenko (2016) reviewed various applications of the eccentric ZLK effect to a broad range of astrophysical systems, such as planetary and exoplanetary systems, stellar systems, and galaxies.
Citation: Naoz, Farr & Rasio 2012
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We noticed that the fast wind stream had He++ beams moving ahead of the protons, low number densities, and low heavy-ion charge state ratios, which are characteristic of coronal-hole-originated solar wind. It should be noted that although the solar wind stream originates from the same coronal hole, the speed and the stream arrival time to 1 au varied from one Carrington rotation to another as was seen in Figure 2. The timing of the coronal-hole-originated solar wind observation aligned well with the PSP orbit. PSP started passing through the leading edge of the high-speed stream, continued as the high-speed stream corotates over the spacecraft, and ended passing the trailing edge of the stream as seen in Figure 3, standing as a perfect classic CIR/SIR event (Cohen et al. 2020). The decrease in crossing width and steepening of the increasing velocity profile at L1 versus at PSP is unlikely to be a product of the evolution of the coronal hole and is most likely a product of the faster wind stream overtaking the slower plasma at the interface during propagation (Burlaga 1974). Analyzing the SIR observed at ACE on 2018 November 4 to the corresponding SIR observed at PSP around 2018 November 15 indicated the enhancement of low-energy (keV) suprathermal ions (Figure 6) at the stream interface as the keV range ions are enhanced by local acceleration. The keV ions were enhanced again after the interface region, but for a short period of time compared to the PSP observation (Figure 8), while high-energy (MeV) suprathermal ions propagate along field lines to the inner heliosphere (Filwett et al. 2019) and are enhanced after the interface region in which articles accelerated at distant shocks dominated (see Desai et al. 2020; Joyce et al. 2021) and lasted longer in time in PSP. Minor enhancements in MeV particles were inconsistently observed over various Carrington rotations at ACE (not shown), but these were often at or just above the instrument noise floor. We cannot make any conclusive inferences as to the evolution of higher-energy ions in the coronal wind from these observations. In the future, as we approach solar maximum, it is likely that there will be more equatorial coronal holes observed during upcoming PSP orbits, which will allow us to compare more in depth the evolution of the solar wind structure at 1 au with PSP.
Citation: Cohen et al. 2020
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The solutions in Sect. 3 show that it is possible for a high X-ray flux in galactic nuclei to alter the carbon ionization balance and reduce the C+ abundance and correspondingly the [C ii] luminosity. In the presence of a high flux of soft X-rays above 1 keV, a condition encountered in many AGNs (Stacey et al. 2010; Ebrero et al. 2009), the abundance of singly ionized carbon is reduced and converted into higher ionization states. This reduction occurs primarily in the hot highly ionized gas that fills most of the galactic central zone, to some degree in the dense ionized skins surrounding molecular clouds, less so in diffuse atomic hydrogen clouds, and very little, if at all, in the dense PDRs at the edge of the CO molecular cores. Furthermore, in galactic nuclei with a high X-ray flux we would expect to see an increase in the dust temperature and infrared luminosity (Voit 1991). Therefore in galactic nuclei, and in particular in AGNs, we expect a reduction in [C ii] emission relative to FIR/IR emission depending on the relative contribution of different ISM components to the [C ii] luminosity. Unfortunately the next most abundant ion, C2+, does not have fine structure FIR emission lines as its 2s2 ground state has spin zero, and C3+, although it has spin angular momentum due to its unpaired electron in the 2S level, does not have a nuclear spin or orbital angular momentum to break the degeneracy of the two electron spin states. Instead to test the effect of X-ray ionization on the carbon balance we would need studies of their UV emission. For example the [C iv] UV resonance lines are detected in extragalactic sources and have been used to trace the formation rate of massive stars (Leitherer & Lamers 1991; Robert et al. 1993). [C iv] UV absorption lines have also been used to study the properties of the Galactic Halo (Savage et al. 2000) where carbon is presumed to be collisionally ionized in hot coronal gas (Gnat & Sternberg 2007).
Citation: Voit 1991
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We note that Pavesi et al. (2019) derived a lower κs parameter log(κs) ≈ −1 for HZ10, by observing the CO(2-1) line, which implies a low star formation efficiency for this source. The conflict between the two results can be explained by the fact that Pavesi et al. (2019) estimated the gas mass by adopting a large CO-to-Mgas conversion factor αCO = 4.5 M⊙ (K km s−1 pc2)−1, a value that is close to the Galactic conversion factor αCO = 4.36 M⊙ (K km s−1 pc2)−1 (Bolatto et al. 2013). Although the Galactic conversion factor is a derived value for Milky Way and normal, star-forming galaxies in the local Universe, it may not be applicable for more extreme environments of starburst galaxies at high-z (see Carilli & Walter 2013 for a review). The conversion factor depends on the physical conditions of the gas in the ISM (temperature, surface density, dynamics, and metallicity), as well as the star formation and associated feedback (Narayanan et al. 2011, 2012; Genzel et al. 2012; Feldmann et al. 2012; Renaud et al. 2019; see, e.g., Bolatto et al. 2013 for a review). It is typically in the range between 0.8 and 4.36 M⊙ (K km s−1 pc2)−1 (see, e.g., Bolatto et al. 2013; Carilli & Walter 2013, and Combes 2018 for reviews). Low metallicities (Z = 0.6 Z⊙ for HZ10) will drive αCO towards values higher than the Galactic value (Narayanan et al. 2012; Genzel et al. 2012; Popping et al. 2014), although αCO spans a broad range of values of αCO ∼ 0.4 − 11 M⊙ (K km s−1 pc2)−1, due to large uncertainties (see, e.g., Fig. 9 of Bolatto et al. 2013). On the other hand, high values of temperature, surface density, and velocity dispersion in a turbulent ISM of starbursts and merging systems will shift αCO towards lower values (Narayanan et al. 2011, 2012; Vallini et al. 2018). HZ10 has an extremely high value of the burstiness parameter log(κs) ∼ 1.4 and high total density of the [C II] emitting gas log(n) ∼ 3.35 cm−3, and it is also a multi-component system (Jones et al. 2017, Carniani et al. 2018a). Thus, for this source we assumed αCO = 0.8 M⊙ (K km s−1 pc2)−1, usually adopted for starburst galaxies (e.g., Downes & Solomon 1998; Bolatto et al. 2013). We obtained log(κs) = 0.53 ± 0.34, which is within the 2σ uncertainties of the 



log

(

κ
s

)

=
1
.

43

−
0.53


+
0.38




$ \log{(\kappa_s)} = 1.43_{-0.53}^{+0.38} $


, estimated exploiting the C III] emission and the Vallini et al. (2020) model.
Citation: Narayanan et al.; 2012
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Citation: Narayanan et al.; 2012
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A substantial fraction of metal-poor stars that have recently evolved off the MS, e.g., giants and planetary nebulae (PNe), have been influenced by binary interactions. The IMF is significantly weighted toward low-mass stars (Bastian et al. 2010; Kroupa et al. 2013), and the MW star formation rate was ≈3 times larger ≈10 Gyr ago than it is now (Governato et al. 2007; De Lucia et al. 2014). Based on the measured IMF and modeled galactic star formation history, we estimate that ≈55% of MW giants and PNe have old, solar-type progenitors (τ* > 7 Gyr, M ≈ 0.8–1.2 





). Such old, low-mass giants tend to be metal-poor (Ratnatunga & Yoss 1991; Carollo et al. 2010; Mackereth et al. 2017). The metallicity trend therefore dramatically affects the properties of low-mass evolved stars. For example, the enhanced close binary fraction of metal-poor solar-type stars substantially strengthens the conclusion that the shaping of PN morphologies is the result of binary interactions (Moe & De Marco 2006; De Marco 2009; Jones & Boffin 2017). Providing further corroboration, Badenes et al. (2015) measured the delay-time distribution of bright PNe in the LMC and discovered two distinct populations of PN progenitors: an old channel (τ* = 5–8 Gyr) deriving from solar-type stars (M ≈ 1.0–1.2 





) and a young channel (35–800 Myr) evolving from late-B/early-A stars (≈2–8 





). According to the measured age–metallicity relation of the LMC (Olszewski et al. 1991; Pagel & Tautvaisiene 1998; Cole et al. 2005; Carrera et al. 2011; Piatti & Geisler 2013), the old, solar-type progenitors are metal-poor ([Fe/H] ≲ −1.0) and hence have a large close binary fraction of Fclose = 40%–50%. The young progenitors have a higher metallicity of [Fe/H] ≈ −0.4 but are sufficiently massive so that they also have a large close binary fraction of Fclose = 40%–60%. Meanwhile, evolved stars with intermediate masses (M ≈ 1.2–2.0 





) in the LMC have intermediate metallicities and therefore a smaller close binary fraction of Fclose ≈ 30%. If PNe derive from interactions in close binaries, then the variations in Fclose with respect to mass and metallicity can explain the observed bimodal mass/age distribution of PN progenitors in the LMC.
Citation: Moe & De Marco 2006
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Citation: Jones & Boffin 2017
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