adsabs/FOCAL huggingface dataset:
Recently Izotov et al. (2016a,b, 2018a,b) have discovered significant emission of Lyman continuum (LyC) ionizing radiation leaking with the escape fractions of 2–76 per cent in a sample of 11 low-z compact active star-forming galaxies (SFGs) observed with the Hubble Space Telescope (HST) in conjunction with the Cosmic Origins Spectrograph (COS). These galaxies hereafter referred to as LyC leakers, possess many properties similar to those of high-redshift galaxies both at z ∼ 2–3 and z$\gtrsim$ 6 such as compact morphology with similar galaxy radii (e.g. Bouwens et al. 2004; Ferguson et al. 2004; Oesch et al. 2010; Ono et al. 2012; Shibuya, Ouchi & Harikane 2015; Curtis-Lake et al. 2016; Paulino-Afonso, Sobral & Ribeiro 2018), strong emission lines with high equivalent widths (EWs; e.g. Schaerer & de Barros 2009; Smit et al. 2014, 2015; Roberts-Borsani et al. 2016; Bowler et al. 2017; Castellano et al. 2017; Fletcher et al. 2019; Bian & Fan 2020; Endsley et al. 2020), similar low stellar masses, low metallicities, and high specific star formation rates (SFRs; e.g. Jaskot & Oey 2013; Nakajima et al. 2013; Stark et al. 2013a; de Barros, Schaerer & Stark 2014; Duncan et al. 2014; González et al. 2014; Nakajima & Ouchi 2014; Becker, Bolton & Lidz 2015; Grazian et al. 2015; Salmon et al. 2015; Huang et al. 2016; Stark 2016; Santini et al. 2017; Stark et al. 2017; Dors et al. 2018), small dust content (e.g. Ouchi et al. 2013; Ota et al. 2014; Maiolino et al. 2015; Schaerer et al. 2015; Watson et al. 2015), and are considered as the main sources of reionization of the Universe after the cosmic ‘Dark Ages’. This makes low-z LyC leakers the best local analogues of reionization galaxies (see e.g. Schaerer et al. 2016; Stark 2016; Ma et al. 2020). Given their proximity, these galaxies represent excellent laboratories for a detailed study of their physical conditions, and the main mechanisms responsible for LyC leakage. Ground-based spectroscopic observations in the visible and near-infrared ranges are necessary for that.
Citation: Ferguson et al. 2004
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A rough idea of the location of the inert Oort cloud can be obtained from previous works. Gladman et al. (2002) showed that the scattering effect by Neptune is significant over long timescales only for perihelion distances below about 45 astronomical units (au). The precise limit actually increases with the semi-major axis (Gallardo et al. 2012), because energy kicks result in larger variations of semi-major axis if the semi-major axis is large. In fact, some observed objects with perihelion beyond 45 au are known to experience scattering (Bannister et al. 2017). In any case, we look here for a rough limit only. The scattering process mostly affects the semi-major axis of small bodies, which diffuses chaotically, while the perihelion distance does not vary much. Later on, Gomes et al. (2005), Gallardo et al. (2012), and Saillenfest et al. (2016, 2017a), showed that the Lidov–Kozai mechanism raised by the giantplanets inside a mean-motion resonance with Neptune is able to raise the perihelion of small bodies beyond 60 au in a few thousands of million years. Contrary to scattering effects, this mechanism induces a variation of perihelion distance and inclination, while the semi-major axis remains at the resonance location. This mechanism, however, is only efficient for semi-major axes smaller than about 500 au. From these studies, one can deduce that the action of the planets is limited to orbits with perihelion distances smaller than about 80 au, and that for perihelion beyond 45 au, the semi-major axis should be smaller than 500 au for the planets to possibly have a substantial effect through mean-motion resonances. As regards the effects of the galactic tides, Fouchard et al. (2017) showed that an object with perihelion in the Jupiter-Saturn region, that is, below 15 au from the Sun, should have a semi-major axis larger than 1600 au for the tides to be able to raise its perihelion beyond 45 au in less than the age of the solar system. In other words, the tides can move its perihelion out of reach of any significant planetary scattering.
Citation: Saillenfest et al.; 2017a
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These signatures are inconsistent with ultra-low frequency waves, which have circular polarization and a period similar to the upstream ion gyroperiod. The waves are also inconsistent with ion Weibel instability, which generates linearly polarized waves. Interaction of reflected ions with incoming solar wind electrons or ions can cause foot instabilities that excite waves in the whistler mode branch. Modified Two Stream Instability (MTSI) due to relative drift between reflected ions and incoming solar wind electrons (fast drift), and incoming solar wind ions and electrons (slow drift) has been frequently considered (Matsukiyo & Scholer 2003; Comişel et al. 2011; Umeda et al. 2012; Marcowith et al. 2016; Wilson et al. 2016; Muschietti & Lembège 2017; Hull et al. 2020). This instability, however, if excited, creates significant ion heating throughout the foot and suppresses the reformation process (Shimada & Hoshino 2005; Matsukiyo & Scholer 2006), rather than creating episodic enhancements that we show in the foot. Furthermore, Gary et al. (1987) indicated that (fast drift) MTSI becomes dominant at low electron beta (βe  0.5), while at higher βe more resonant electrons stabilize this instability through increased electron Landau damping. Electron data for the time period we discussed in this paper show βe ≥ 1.2, and therefore fast drift mode MTSI is most likely not significant. The slow drift mode of MTSI could be a more viable candidate at high β plasmas. Wave properties around 1.6 Hz in the middle interval (purple segment) of Figure 6, indicate that the wave is in propagation toward the ramp (





 = −0.66, −0.71, 0.22) with Vph‐sc = 34 km s−1 and λwave = 21.4 km  ∼ 12λe, where λe is the upstream electron inertial length. The plasma rest-frame frequency of the wave is about 8 Hz  ∼ 2.5flh. Since these characteristics are somewhat consistent with model predictions for drift mode of MTSI (Muschietti & Lembège 2017), we do not rule out the possibility of some waves at certain frequencies and during some intervals being generated by the slow drift mode of MTSI.
Citation: Muschietti & Lembège 2017
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Citation: Muschietti & Lembège 2017
Citation function: Similarities; Similarities
A lower CB frequency in GCs than in the field has already been noticed by Rucinski (2000). It results most likely from a very low binary frequency of ~0.1 in GCs (but with a significant scatter among different clusters), compared to 0.5 in the field (Milone et al. 2012), together with a low percentage of binaries in the contact phase at a given age, as shown above. Any statistical comparison of models with observations is therefore extremely uncertain because we have so few observational data. Nevertheless, we can compare the predicted to observed fraction of BSs among all CBs in a GC. We used data for two clusters observed by Kaluzny and his collaborators with the highest number of CBs (except for ω Centauri). In M4, Kaluzny et al. (2013a) detected nine CBs with one BS among them, and in 47 Tuc, Kaluzny et al. (2013b) identified 15 CBs or NCBs, of which six are BS. The resulting fractions are 0.1 and 0.4, respectively. We did not take into account ω Cen with the richest population of CBs (Kaluzny et al. 2004) because it is highly atypical and the accurate number of member CBs is not well known, although the approximate data indicate a similarly high ratio as in case of 47 Tuc. The predicted fractions range from 0.08 to 0.2 for the three considered cluster ages. As we see, they are close to the lower observed value but are at odds with the higher value. This may indicate some deficiencies of the CCBM, for instance, too short initial cut-off period, a too high AML rate at short periods, or a too low mass-transfer rate. With the lower AML rate and/or higher mass transfer rate, a binary stays longer in contact and can reach a lower mass ratio before merging. Many of the short-period CBs have a rather high mass ratio of about 0.7−0.8 at the time of merging, whereas field CBs with P ≲ 0.3 d center around a value of 0.5 (Rucinski 2010; Stȩpień & Gazeas 2012). The lower mass ratio means more mass transferred to the gainer, hence its higher mass and higher position on the HRD, that is, a higher probability for entering the BS region. There may also still be another explanation of the discrepancy: strong fluctuations of this fraction among different clusters suggest a nonuniform BSs formation rate in some clusters, with individual bursts occurring in the recent past (see below). Regardless of the reason, it is apparent that profound differences occur among GCs, which makes a comparison of theoretical predictions with individual clusters uncertain.
Citation: Kaluzny et al. (2013a)
Citation function: Uses; Background; Background
HD 169142 is a very young Herbig Ae-Be star with a mass of 1.65–2 M⊙ and an age of 5–11 Myr (Blondel & Djie 2006; Manoj et al. 2007) that is surrounded by a gas-rich disk (i = 13°; Raman et al. 2006; PA = 5°; Fedele et al. 2017) that is seen almost face-on. The parallax is 8.77 ± 0.06 mas (Gaia DR2 2018). Disk structures dominate the inner regions around HD 169142 (see, e.g., Ligi et al. 2018). Figure 1 shows the view obtained from polarimetric observations: the left panel shows the QΦ image in the J band obtained by Pohl et al. (2017) using SPHERE on a linear scale, and the two rings are clearly visible. The right panel shows a pseudo-ADI image of the inner regions obtained by differentiating the QΦ image (see Ligi et al. 2018, for more details). Biller et al. (2014) and Reggiani et al. (2014) discussed the possible presence of a point source candidate at small separation (0.2 arcsec from the star). However, the analysis by Ligi et al. (2018) based on SPHERE data does not support or refute these claims; in particular, they suggested that the candidate identified by Biller et al. (2014) might be a disk feature rather than a planet. Polarimetricimages with the adaptive optics system NACO at the Very Large Telescope (VLT; Quanz et al. 2013b), SPHERE (Pohl et al. 2017; Bertrang et al. 2018) and GPI (Monnier et al. 2017) show a gap at around 36 au, with an outer ring at a separation >40 au from the star. This agrees very well with the position of the rings obtained from ALMA data (Fedele et al. 2017); similar results were obtained from VLA data (Osorio et al. 2014; Macías et al. 2017). We summarize this information about the disk structure in Table 1 and call the ring at 0.17–0.28 arcsec from the star Ring 1 and the ring at 0.48–0.64 arcsec Ring 2. We remark that in addition to these two rings, both the spectral energy distribution (Wagner et al. 2015) and interferometric observations (Lazareff et al. 2017; Chen et al. 2018) show an inner disk at a separation smaller than 3 au. This inner disk isunresolved from the star in high-contrast images and consistent with ongoing accretion from it onto the young central star.While the cavities between the rings seem devoid of small dust, some gas is present there (Osorio et al. 2014; Macías et al. 2017; Fedele et al. 2017). Fedele et al. (2017) and Bertrang et al. (2018) have suggested the possibility that the gap between Rings 1 and 2 is caused by a planet with a mass slightly higher than that of Jupiter. However, this planet has not yet been observed, possibly because it is at the limit of or beyond current capabilities of high-contrast imagers. On the other hand, Bertrang et al. (2018) found a radial gap in Ring 1 at PA ~ 50° that might correspond to a similar radial gap found by Quanz et al. (2013b) at PA ~ 80°. The authors noted that if this correspondence were real, then this gap might be caused by a planet at about 0.14 arcsec from the star. So far, this planet has not been unambiguously detected either.
Citation: Pohl et al. (2017)
Citation function: Uses; Uses
Citation: Pohl et al. 2017
Citation function: Uses; Uses; Uses
As part of this study, we investigate the 8 μm emission within our sample of galaxies using archival data from IRAC on the Spitzer Space Telescope (SST). The rest-frame 8 μm emission from galaxies has historically been used as a monochromatic SFR indicator. This is justified by the fact that polycyclic aromatic hydrocarbons (PAHs) and small dust grains can be heated by the photodissociation regions (PDRs) that surround actively star-forming regions and the subsequent emission from PAHs is brightest at about 8 μm (Draine & Li 2007; Smith et al. 2007; Kennicutt et al. 2009; Elbaz et al. 2011). However, there is an important caveat; PAHs are destroyed by ionizing radiation from newly formed stars (Helou et al. 2004; Povich et al. 2007; Bendo et al. 2008; Relaño & Kennicutt 2009), leading to a deficit in the 8 μm luminosity in galaxies with low metallicity, whereas PAHs are less well shielded by metals (Engelbracht et al. 2005; Calzetti et al. 2007; Smith et al. 2007; Cook et al. 2014; Shivaei et al. 2017). It has also been argued that metals can act as catalysts for the formation and growth of PAHs, leading to smaller average sizes in low-metallicity environments (Sandstrom et al.2012). More recently, Lin et al. (2020) find an anticorrelation between the dust-only 8 μm luminosity and the age of young stellar clusters, suggesting the 8 μm luminosity decreases with increasing age of the stellar population. The existence of a strong interstellar radiation field is also found to suppress the emission from PAHs, independent of metallicity (Madden et al. 2006; Gordon et al. 2008; Lebouteiller et al. 2011; Shivaei et al. 2017; Binder & Povich 2018). Adding to the complexity, spatially resolved studies have shown that a significant amount of 8 μm emission is associated with the cold, diffuse ISM, which suggests an important heating source other than recent (100 Myr) star formation (Bendo et al. 2008; Calapa et al. 2014; Lu et al. 2014). These results suggest that the luminosity at 8 μm is not a straightforward indicator of the SFR. In this work, we further investigate these issues by exploring how metallicity gradients within our sample affect the observed “red” side versus “blue” side IR color–color correlations—where we derive the “blue” side color from the flux density ratios at 8 and 24 μm.
Citation: Calzetti et al. 2007
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Post-main-sequence mergers. Mergers of evolved stars with their MS companions originate from wider systems than main-sequence mergers (log P ≳ 0.7). We thus roughly estimate 




f

P
,
i


∼

IPF

0.7


3.0


∼
0.66


$ f_{P , \mathrm{i}} \sim \mathrm{IPF}_{0.7}^{3.0} \sim 0.66 $


, using log10P = 3 as the maximum initial period for binary interaction (e.g., Claeys et al. 2011; Yoon et al. 2017). For merging to occur, a CEE phase needs to be initiated. Extreme mass ratio systems are prone to unstable mass transfer and thus to CEE, and although the exact value of the boundary is not well-constrained, we assume that this occurs for evolved donors in systems with q ≲ 0.4, following previous works such as Wellstein et al. (2001), Hurley et al. (2002) and de Mink et al. (2013). In principle, CEE can alternatively lead to the ejection of the envelope, but in this simple estimate we assume that all these systems eventually merge, leaving some hydrogen-rich layers on the surface of the formed star. This is consistent in most of the cases with the findings from our computational simulations for our standard assumption of using the entire orbital energy change to eject the envelope (αCEE = 1, as we will introduce and discuss in Sect. 4.1). Thus, 




f

q
,
i


∼

IQF

0.1


0.4


∼
0.33


$ f_{q , \mathrm{i}} \sim \mathrm{IQF}_{0.1}^{0.4} \sim 0.33 $


. The projection of the initial parameter space for this channel at the log P − q plane is depicted in magenta on the left panel of Fig. 2. The donor star in such systems originate from a roughly similar part in M1 space as main-sequence mergers (




f


M
1

,i

~

IMF
7

25

~1.25

$ f_{M_1 , {\rm i}} \sim {\rm IMF}_7^{25} \sim 1.25 $



). In contrast, binaries of less extreme mass ratio are able to either follow a phase of stable mass transfer onto the secondary star or are assumed to always survive a CEE by ejecting the envelope, avoiding a merger. Such donors are expected to eventually produce a hydrogen-poor, stripped-envelope SN, not contributing to the SN II population that we focus on in this study. We find XpostMS + MS mergers ∼ C × 0.136 ∼ 15%.
Citation: de Mink et al. (2013)
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We further restricted our sample to well-defined N2H+ (1–0) spectra that we used to estimate the gas velocity dispersion. The N2H+ (1–0) emission of each clump was evaluated from the MALT90 data cubes by averaging the spectrum across all the pixels within one MALT90 beam, ≃ 38 arcsec. We assumed that all the N2H+ emission comes from the clumps, and we estimate the filling factor from the comparison of the radius of each Hi-GAL clump with the radius of a region equal to the MALT90 beam (Fig. 1). There is a strong correlation between these two quantities, and the size of the Hi-GAL clumps is systematically smaller than the radius estimated from the MALT90 beam for a factor of 0.64, on average. We assumed an average filling factor of 0.64 for the entire sample. The MALT90 data cubes are given in antenna temperature $T_{\rm A}^{*}$ and they have been converted to the main beam temperature $T_{\rm MB}=T_{\rm A}^{*}/\eta _{\rm MB}$, assuming a mean beam efficiency ηMB = 0.49 (Miettinen 2014). The properties of each N2H+ (1–0) averaged spectrum have been extracted in idl using a hyperfine fitting routine and the mpfitfun algorithm (Markwardt 2009), after smoothing the data to a spectral resolution of 0.3 km s−1 to enhance the signal-to-noise (S/N) ratio. We excluded all clumps with a S/N ratio below 5, where the rms in each smoothed data cube has been measured in a 100 km s−1 wide spectral window near the N2H+ emission. We further excluded clumps for which the fit converged but the spectrum was affected by spikes and/or by multiple components along the line of sight. Using these criteria, we obtained 308 clumps. We completed our selection by excluding all clumps without a clear distance assignation, in particular without a well-defined resolution of the near–far distance ambiguity. First, we have refined the kinematic distances in the Elia et al. (2017) catalogue (and the quantities that depend on them) with the newest set of distances developed for the Hi-GAL survey under the VIALACTEA project (Mege et al., in preparation). The method used by Elia et al. (2017) was the same as in Russeil et al. (2011): the brightest emission lines in the 12CO or 13CO spectra along the line of sight of each source are used to estimate the velocities of the local standard of rest and converted into heliocentric distances using the Brand & Blitz (1993) rotation curve. The distances in Mege et al. (in preparation) have been determined using a similar approach, but including all the recent surveys of the Galactic plane to trace structures along the line of sight, and using the more recent Reid et al. (2009) rotation curve. Then, in order to identify only clumps with a well-defined distance estimation, we have compared the distances assigned to our 308 clumps with the distances of the MALT90 sample estimated in Whitaker et al. (2017) and of the ATLASGAL sources published in Urquhart et al. (2018). We excluded from the sample all sources with a difference in the distance estimation larger than 20 per cent among the three surveys.
Citation: Miettinen 2014
Citation function: Uses
In the present work, we neglect the processes that generate the halo, but this topic deserves some review. Notably, the apparent growth of the halo at the expense of the anti-sunward suprathermal ‘strahl’ population may imply that the halo is locally formed in the inner heliosphere by scattered strahl electrons (e.g. Maksimovic et al. 2005; Štverák et al. 2009). This has led to significant theoretical development, focused on the resonant interaction of electrons with the whistler and fast-magnetosonic whistler (FM/W) modes (e.g. Vocks et al. 2005; Saito & Gary 2007; Vasko et al. 2019; Verscharen et al. 2019b; Micera et al. 2021; Zenteno-Quinteros, Viñas & Moya 2021; Tang, Zank & Kolobov 2022; Vo, Lysak & Cattell 2022). Observations have struggled to confirm these theories. Notably, whistlers are practically absent (occurrence rate 0.1 per cent) during PSP perihelion passes (Cattell et al. 2022). Additionally, the eVDFs sampled by Helios and PSP are stable with respect to the oblique FM/W mode (Jeong et al. 2022a). Theoretical calculations show that at r ≲ 1 AU, the strahl is stable to whistler fluctuations (Horaites et al. 2018b; Schroeder et al. 2021) and should be unaffected by whistler turbulence in the inner heliosphere (Boldyrev & Horaites 2019). High-resolution measurements of the strahl at 1 AU confirm that ‘anomalous diffusion’, e.g. from whistler waves, is not required to explain the strahl angular widths at resolvable energies ≲300 eV Horaites et al. (2018a), Horaites, Boldyrev & Medvedev (2019). Similar results were found from simulations at distances r ≲ 20RS (Jeong et al. 2022b), which showed that near the corona the strahl is adequately described by a combination of Coulomb collisions and expansion effects. This all suggests that a mechanism besides local wave particle scattering may account for the halo’s presence in the inner heliosphere. Such theories have been proposed (e.g. Leubner 2004; Lichko et al. 2017; Che et al. 2019; Horaites et al. 2019; Scudder 2019), though no consensus has emerged.
Citation: Zenteno-Quinteros, Viñas & Moya 2021
Citation function: Background; Differences
The currently accepted cold dark matter dominated model with the cosmological constant (ÎCDM) predicts that structures in our Universe assemble hierarchically, with more massive systems forming later through accretion and mergers of smaller, self-bound dark-matter halos (e.g., Tormen 1997; Moore etÂ al. 1999; Klypin etÂ al. 1999; Springel etÂ al. 2001). In N-body cosmological simulations, dark matter halos of all masses converge to a roughly âuniversalâ and cuspy density profile that steepens with radius, the so-called NavarroâFrenkâWhite profile (NFW profile; Navarro etÂ al. 1996, 1997). Moreover, the degree of central concentration of a halo depends on its formation epoch and hence on its total mass (e.g., Wechsler etÂ al. 2002; Zhao etÂ al. 2003). Within this scenario, early virialized objects are compact when they get accreted into a larger halo. Such objects are usually referred to as subhalos or substructures of their host and, as they orbit within the host potential well, they are strongly affected by tidal forces and dynamical friction, causing mass, angular momentum, and energy loss (e.g., Ghigna etÂ al. 1998; Tormen etÂ al. 1998; De Lucia etÂ al. 2004; Gao etÂ al. 2004). In the ÎCDM framework, more massive halos are predicted to have a larger fraction of mass in subhalos than lower mass halos because in the former there has been less time for tidal destruction to take place (e.g., Gao etÂ al. 2004; Contini etÂ al. 2012). On galaxy cluster scales, observational tests of these predictions have been attempted in some previous works (e.g., Natarajan etÂ al. 2007, 2009), but highly accurate analyses are becoming possible only now, thanks to the substantially improved quality of the available photometric and spectroscopic data. From an observational point of view, more investigations are still required to fully answer key questions on the formation and evolution of subhalos. How much mass of subhalos is stripped as they fall into the host potential? How many subhalos survive as bound objects? What are the spatial and mass distributions of the subhalos?
Citation: Zhao etÂ al. 2003
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