adsabs/FOCAL huggingface dataset:
The short-term (seconds to hours) temporal properties of ULXs are quite different from those observed in Galactic X-ray binary systems and are not related to specific spectral states (e.g. Heil et al. 2009): even if the amplitude of short-term variability is larger in soft ultraluminous sources than in hard ones (Sutton et al. 2013), the variability appears to be poorly predictable, i.e. it is not found in all sources with a soft spectrum and, when detected in a source, it is not necessarily found in all the observations of that ULX. Such variability was explained with the existence of optically thick and non-uniform winds, radiatively ejected by the accretion disc, which stochastically intersect our line of sight (e.g. Middleton, Sutton & Roberts 2011; Middleton et al. 2015a). Observational evidences of these winds come from the detection of blue-shifted absorption lines in the high-quality grating spectra of some ULXs (e.g. Pinto, Middleton & Fabian 2016; Kosec et al. 2018) and by the nebulae observed in the radio (e.g. Cseh et al. 2012), optical (e.g. Pakull, Soria & Motch 2010) or X-ray band (Belfiore et al. 2020). Long-term flux variability (days to months) is also observed in most of the monitored ULXs. Such variability, in some cases, can be as high as several orders of magnitudes, implying that these sources are transients (detected at least once in the ultraluminous state, Lx > 1039 erg s−1, and either once at a significantly smaller luminosity or undetected below the instrumental sensitivity; e.g. Soria et al. 2012; Pintore et al. 2018a). Often the pulsating ULXs (PULXs) show large flux variations, implying they can be considered as transient ULXs. A bimodal flux distribution is sometimes observed in the long-term light curves of PULXs (e.g. Walton et al. 2015a; Motch et al. 2014). A possible explanation for this long-term behaviour may be the propeller effect: when the magnetospheric radius of the NS becomes larger than the corotation radius of the accreting matter in the disc, a centrifugal barrier prevents the accretion on to the NS, with a corresponding decrease in the X-ray flux (Illarionov & Sunyaev 1975; Tsygankov et al. 2016; Grebenev 2017). The propeller effect can act only if the accretor is an NS with a strong magnetic field, so a bimodal flux distribution can be used to identify candidate PULXs, even when pulsations are not detected (e.g. Earnshaw, Roberts & Sathyaprakash 2018; Song et al. 2020). Furthermore, long-term monitored ULXs revealed possible (super-)orbital variability (e.g. Foster, Charles & Holley-Bockelmann 2010; An, Lu & Wang 2016; Fürst et al. 2018), the origin of which is still a matter of debate. Super-orbital periods are also observed in the light curves of the PULXs, with periods of tens to hundreds days (e.g. Walton et al. 2016; Brightman et al. 2019, 2020). The detections of such long-term periodicities were only possible thanks to the flexibility and performance of the Neil Gehrels Swift Observatory (hereafter Swift; Gehrels et al. 2004).
Citation: Pinto, Middleton & Fabian 2016
Citation function: Background
X Persei (4U 0352 + 309) is a binary stellar system at a distance of ∼0.95 kpc (Telting et al. 1998) and consists of a slowly spinning neutron star (White, Mason & Sanford 1977, Pspin ∼ 837 s) that is accreting matter from its Be-star companion (Lyubimkov et al. 1997). The source is peculiar in that it is a persistent source and does not show Type-I outbursts as are commonly observed in other Be-/X-ray binaries. But X-ray flares and variability in the X-ray light curve of the source have been observed (La Palombara & Mereghetti 2007; Lutovinov, Tsygankov & Chernyakova 2012). The orbital period of the binary system is long, $P_{\text{orb}}\sim 250 \rm {d}$, but the orbital eccentricity is only e ∼ 0.11 (Delgado-Martí et al. 2001). The long orbital period and the low X-ray luminosity of the source support the assumption that the neutron star in this binary is accreting quasi-spherically by capturing matter from the stellar wind of the Be-star (Shakura et al. 2012). The X-ray spectrum of X Persei is also unlike most accreting X-ray pulsars, the latter can usually be modelled with a power law and a cutoff in the range 10–30 keV (White, Swank & Holt 1983). Instead, X Persei has a peculiarly hard X-ray spectrum, making the source detectable at energies higher than 100 keV (Lutovinov et al. 2012). Thus, a single-component spectral model is insufficient to model the source continuum spectrum, making it necessary to include a second high-energy component to explain the emission at harder X-rays. Several authors have used several different models to explain the observed spectrum, of which we note the following works. The RXTE observations were modelled using a blackbody component at lower energies and a power-law component with an exponential cutoff at higher energies (Coburn et al. 2001). A cyclotron resonance scattering feature (CRSF) was also required to account for an absorption-like feature seen around 30 keV. The spectrum obtained from BeppoSAX observation was modelled with two power-law components with exponential cutoffs at low and high energies and did not require any CRSF (Di Salvo et al. 1998). The INTEGRAL observations used thermal Comptonization and bulk motion Comptonization to account for the observed spectrum at the low and high energies, respectively (Doroshenko et al. 2012). The INTEGRAL observations also did not require any CRSF to improve the spectral model. Thus, the presence of CRSF noted in the RXTE observations was not detected in either the BeppoSAX or the INTEGRAL observation and hence needs to be verified. There is also no clear agreement on the unusual nature of the X-ray spectrum with respect to the continuum model.
Citation: Lutovinov et al. 2012
Citation function: Compare/Contrast; Motivation
The lower panel of Figure 9 shows the SFR density against the H2 mass density for each pixel (black dots) and the average value found within each of the analyzed region as colored symbols. Red dashed lines are fixed depletion times, while the blue dashed line is the one derived from the 30 nearby disk galaxies of the HERACLES survey by Bigiel et al. (2011) at 1 kpc scale resolution. The red dots, corresponding to regions D1, D2, and D3 located within the disk, lie below the local galaxies relation (in blue), i.e., they show low SFE. This confirms that this galaxy is forming new stars in the disk at a very low rate (see also Vulcani et al. 2018), given that the measured values of 





 are those typically found in galaxy bulges (Fisher et al. 2013). Lower values of αCO (5–10 times lower than the Milky Way value adopted here) have been found in the central part of some HERACLES galaxies (Sandstrom et al. 2013), but only in the central kiloparsec, while regions D1, D2, and D3 span a larger extent of the galaxy disk. In particular, the central kiloparsec is dominated by AGN-like line ratios (Poggianti et al. 2017a; Radovich et al. 2019), according to different indicators (Poggianti et al. 2019a), and we therefore excluded it from the calculation. Our data confirm that local conditions play an important role in determining the star formation process, as already shown in the nearby galaxy M51 (Bigiel et al. 2016), and in 29 other nearby galaxies where dense gas tracers were available (Usero et al. 2015). Ram pressure then works in JW100 by enhancing the gas density in the disk and at the same time suppressing the global SFR, resulting in long depletion times both in the tail and in the galaxy disk (as already suggested in Moretti et al. 2018a). Whether this is accompanied also by an enhancement of the molecular gas fraction, as in Nehlig et al. (2016), cannot be stated yet. Ongoing H i observations with MeerKAT will shed light on this issue. It is worth noticing, though, that in the GASP jellyfish galaxy JO206, the H i depletion time turned out to be shorter than expected (Ramatsoku et al. 2019), showing the opposite trend with respect to the molecular gas in JW100.
Citation: Vulcani et al. 2018
Citation function: Similarities
Since polycyclic aromatic hydrocarbons (PAHs) were first hypothesized to be carriers of the unidentified infrared bands (Léger & Puget 1984; Allamandola et al. 1985), a great deal of effort has been devoted to understanding the chemical processes leading to the formation of these molecular species (see e.g. Joblin & Cernicharo 2018). Circumstellar envelopes around carbon-rich asymptotic giant branch (AGB) stars have been suggested as the factories of PAHs (Cherchneff et al. 1992). The detection of benzene in the carbon-rich protoplanetary nebula CRL 618 (Cernicharo et al. 2001) suggests a bottom-up approach in which the small hydrocarbons that formed during the AGB phase, such as C2H2 and C2H4, interact with the ultraviolet (UV) radiation produced by the star in its evolution to the white dwarf phase (Woods et al. 2002; Cernicharo 2004). Other hypotheses involve the processing of dust grains around evolved stars, either through UV photons (Pilleri et al. 2015) or by chemical processes (Martínez et al. 2020). Hence, it has been surprising to see that cyanide derivatives of PAHs have been found in the cold pre-stellar core TMC-1, which is well protected against UV radiation (McGuire et al. 2018, 2021). It is unlikely that these PAH cyanides arise from a reservoir of PAHs that has existed since the early stages of the cloud since these relatively small PAHs would not have survived the diffuse cloud stage. Although McGuire et al. (2021) propose a reasonable chemical network starting with the phenyl radical that could explain the observed abundances of cyanonaphthalene, the chemical routes leading to benzene itself are still unclear. An in situ formation mechanism for benzene must involve abundant hydrocarbons containing from two to four carbon atoms. Moreover, some of these species have to permit an easy cyclization in two to three steps to have an efficient yield of benzene or phenyl radical. The propargyl radical (CH2CCH) was recently found in TMC-1 by Agúndez et al. (2021) with an abundance close to 10−8 relative to H2. In addition, complex hydrocarbons such as vinyl and allenyl acetylene have also been observed in very large abundances (Cernicharo et al. 2021a,b). These hydrocarbons may hold the key to the formation of initial aromatic rings, from which larger PAHs can grow.
Citation: Woods et al. 2002
Citation function: Background; Background
For the sources with the full JHKs photometric information (309 objects out of 507 found in the 2013 HKs catalogue), we first check the source's position in the CCD diagram. In Fig. 4(c), the solid black line represents evolutionary models, the dashed black and grey lines the locus of T-Tauri stars and the corresponding uncertainties (Meyer et al. 1997), whereas the dash–dotted rectangle represents the locus of Herbig AeBe stars (Hernández et al. 2005). The dotted red lines are the reddening vectors (Cardelli et al. 1989), encompassing the regions where the colours are consistent with reddened evolutionary models (region A), CTTSs (region B) and Herbig AeBe stars (region C). If the star falls in region A of the CCD, its extinction and corresponding mass are derived by de-reddening its photometry to the 1 Myr isochrone in the J, (J − H) CMD. In region B, the extinction is derived by de-reddening the colours to the CTTS locus (Meyer et al. 1997). We note that the objects in region A could also be part of the CTTSs; however, the derived AV differences are typically smaller than the derived uncertainties, and for simplicity we decide to only consider the evolutionary models. The derived extinction is then used to convert the J-band photometry to the absolute J-band magnitude, which is in turn compared to the models to derive the mass. For the stars in region B, in addition to the interstellar extinction, we also correct for an excess due to the circumstellar disc or envelope, chosen randomly in the interval 0–0.7 mag for the J band (Cieza et al. 2005). In the case that the derived mass is larger than 2 M⊙ (upper mass limit for T-Tauri stars), the procedure is repeated by de-reddening the colours to the Herbig AeBe locus (Hernández et al. 2005). The intrinsic colour in this case is a randomly chosen value within the dash–dotted box shown in Fig. 4(c), falling along the reddening line. The same procedure is performed if the object falls within region C. Finally, if the object falls to the left of region A, or to the right of region C, the extinction and mass cannot be derived.
Citation: Cardelli et al. 1989
Citation function: Uses; Uses
A great deal of this progress is attributed to reverberation mapping (RM) campaigns since the 1980s, with the underlying principle proposed by Bahcall et al. (1972) and Blandford & McKee (1982). Photons of emission lines from structured ionized gas trace different paths to observers, leading to time lags (τ) of the emission lines with respect to the ionizing photons. Such RM campaigns focusing on broad Balmer lines had detected the anticipated lags in a number of Seyfert galaxies (e.g., Peterson 1993; Peterson et al. 1998; Bentz et al. 2013; Barth et al. 2015; U et al. 2022) and quasars (e.g., Kaspi et al. 2000; Du et al. 2014, 2018a; Shen et al. 2019) over the past few decades. Along with the growing investment of observing resources and the development of analytical algorithms, the general geometry and kinematics of BLRs in some AGNs have been revealed by velocity-resolved delay analysis (e.g., Bentz et al. 2010; Denney et al. 2010; Grier et al. 2012, 2013; Du et al. 2016a; U et al. 2022), velocity-delay maps (e.g., Xiao et al. 2018; Horne et al. 2021), or dynamical modeling (e.g., Bottorff et al. 1997; Pancoast et al. 2014; Li et al. 2018; Williams et al. 2020). Many resolved BLRs have a disk-like geometry of their Hβ line region (the other have inflow or outflow, or a kind of mixture of the three configurations)1. Moreover, the repeat RM observations of the same emission line and the RM results of the emission lines with different ionization in a few AGNs (e.g., NGC 5548, 3C 390.3, NGC 3783, NGC 7469, Mrk 817 etc.) approximately demonstrated the relation VFWHM ∝ τ−1/2, demonstrating evidence for potential of SMBHs (e.g., Peterson & Wandel 2000; Peterson et al. 2004; Lu et al. 2021), where VFWHM is the full-width-half-maximum (FWHM) of the emission lines. Considering the disk-like geometry of BLRs in some AGNs, this relation probably indicates a nearly Keplerian rotation of the disk BLRs. More recently, the GRAVITY instrument mounted in Very Large Telescope Interferometer (VLTI) spatially resolved the BLRs in several AGNs (e.g., 3C 273, NGC 3783, IRAS 09149 by GRAVITY Collaboration 2018, 2020b, 2021, respectively) and also found that their BLRs are approximately characterized by Keplerian rotating disks.
Citation: Du et al. 2014
Citation function: Background; Background; Background
We combine theoretical prescriptions with realistic simulations of asteroseismic observations to construct the first magnetic detection method based on the period spacing of mixed modes. This study opens the way for an extensive search of magnetic fields inside evolved stars with solar-like oscillations. Due to the known effect of magnetic fields on mixed-mode frequencies from Bugnet et al. (2021), the magnetic signature on local measurements of the period spacing of g-modes investigated here should not get mistaken for another known physical process also leading to frequency shifts (e.g., latitudinal differential rotation, centrifugal effects, glitches, near-degeneracy effects). Such a detection of magnetic-field signatures from the measurement of the period spacing of g-modes can then be inverted to get an estimate of the radial magnetic-field strength inside RGs following Mathis et al. (2021), as it is already possible for the internal rotation (e.g., Deheuvels et al. 2017, 2020). Our methodology is well-suited for stars observed with low-inclination angles or observed pole-on, as it is based on azimuthal order m = 0 components to avoid additional signatures in the LS periodogram resulting from the distinct periodicity of m = ±1 modes. As long as the magnetic-field amplitude remains moderate, the method presented here can be applied even if the field is inclined inside the star (see Appendix F). From the statistics on the amplitudes of oscillations on the RGB (e.g. Stello et al. 2016a), intermediate-mass stars seems more likely to host strong magnetic fields than low-mass stars during the RGB. There are a few stars in the study of Gehan et al. (2020) that are observed with low-inclination angles, with a mass above 1.3 solar masses, and that present a rather low value of ΔΠ1 from the study of Vrard et al. (2016). The detection of magnetic fields inside these stars is out of the scope of this paper and its focus on methodology, but we are able to identify these stars as the best sample for the search of buried magnetic fields. The observational search is made possible thanks to the theoretical progress made in this study. Such a detection would be game-changing for improving the understanding of all stars of low and intermediate mass, as they are described as non-magnetic bodies in current stellar evolution models. The presence of strong internal magnetic fields inside the radiative interior of stars along their evolution would modify stellar age estimates and would thus affect many astrophysical areas – from planet habitability2 studies to galacto-archeology (Rauer et al. 2014; Chaplin et al. 2020).
Citation: Deheuvels et al.; 2020
Citation function: Uses
As given in equations (1)–(3), ϕ(r) and i(r) are needed when deriving the LOS model velocity at a projected sky position (x, y). Specifically, equations (2) and (3) imply that
(4)
\begin{eqnarray}
r &amp;=&amp; \Biggl [\biggl \lbrace -(x-x_{C}) \, \sin \phi + (y-y_{C}) \, \cos \phi \biggr \rbrace ^2 \nonumber \\
&amp;&amp;+\,\biggl \lbrace \frac{(x-x_{C}) \, \cos \phi \ + (y-y_{C}) \, \sin \phi }{\cos i} \biggr \rbrace ^2\Biggr ]^{\!1/2}.
\end{eqnarray}
If ϕ and i are independent of r, the latter can be directly derived from equation (4). However, if not, adequate functional forms that provide a sufficient approximation to the radial variations of ϕ and i should be assumed. As discussed earlier, kinematic ϕ and i can vary with galaxy radius due to dynamical structures in galaxies including lopsidedness, warps, bars, spiral arms and non-circular motions. The combined effect of such structures tends to result in random variations of ϕ and i which are not necessarily described by any specific functional form. To remove any unphysical discontinuities of ϕ and i and regularize their radial variations, we use the basis spline (de Boor 1978), also called the ‘B-spline’. This is a piecewise radial polynomial function of degree n where the order n is less than the number of rings in the tilted-ring model. The radial extent of the galaxy is broken up into some number of intervals where each interval has two endpoints, called ‘breakpoints’. For continuity and smoothness, these breakpoints are converted to ‘knots’ which constitute a knot vector
(5)
\begin{equation}
\displaystyle {\boldsymbol t} = \{t_{0}, t_{1}, \ldots, t_{n+k-1}\},
\end{equation}
where n is the number of basis splines of order k.
(6)
\begin{eqnarray}
\displaystyle B_{m,1}(r) &amp;=&amp; \left\lbrace \begin{array}{rl}1 &amp; \qquad t_{m} \le r < t_{m+1} \\
0 &amp; \qquad \rm {otherwise} \end{array} \right.
\end{eqnarray}
(7)
\begin{eqnarray}
B_{m,k}(r) &amp;=&amp; \frac{r-t_{m}}{t_{m+k-1} - t_{m}}B_{m,k-1}(r) \nonumber \\
&amp;&amp;+\,\,\frac{t_{m+k} - r}{t_{m+k} - t_{m+1}} B_{m+1,k-1}(r),
\end{eqnarray}
where m = 0, 1, …, n − 1. Constant, linear, quadratic
and cubic B-splines are given by k = 1, 2, 3 and 4, respectively. The models of ϕ and i used in the new algorithm are given by expanding the B-spline functions as follows:
(8)
\begin{equation}
\phi (r) = \sum _{l=1}^{U} c^{\phi }_{u}B^{\phi }_{l, k}(r),
\end{equation}
(9)
\begin{equation}
i(r) = \sum _{m=1}^{V} c^{i}_{v}B^{i}_{m, k}(r),
\end{equation}
where U and V are the numbers of B-splines, and $c^{\phi }_{m}$ and $c^{i}_{m}$ are the coefficients of the B-splines for ϕ and i, respectively. Similarly, the expansion velocity, vEXP, can also be modelled by the expansion of W B-spline functions,
(10)
\begin{equation}
v_{\rm EXP}(r) = \sum _{n=1}^{W} c^{v_{\rm EXP}}_{w}B^{v_{\rm EXP}}_{n,k}(r).
\end{equation}
Citation: de Boor 1978
Citation function: Uses; Uses
From the plot, we can see that there are some outliers in the sample. We will discuss these systems specifically in the following. For the system with an open circle, it is more likely to have a WD mass larger than 0.48 $\rm M_{\odot }$ and could have a CO WD rather than a He WD. Among these observed data, there are three systems with small orbital periods (log(Porb/days)  −1.2) (open squares). For the two systems with orbital periods ($\rm {log} (P/\rm {days})$) around –1.6, their sdB masses roughly equal to 0.3 $\rm M_{\odot }$ (Burdge et al. 2020). Because only these sdB progenitors with non-degenerate cores can ignite the helium when their core masses are around 0.3 $\rm M_{\odot }$ (Han et al. 2002, 2003). Their progenitor masses of sdB stars are likely to be larger than $2.0\,\,\rm M_{\odot }$. For these progenitors, they have a larger binding energy of envelopes leading to smaller orbital periods after CE ejection. For the system PTF1 J0823, Kupfer et al. (2017) claimed that the progenitor of the sdB star in this system is likely to have a non-degenerate core. For the three systems with minimum WD masses smaller than $0.20\,\,\rm M_{\odot }$, it is unclear how the systems formed if the WD masses are confirmed to be less than 0.2 $\rm M_{\odot }$. In principle, stellar objects with such a low mass could be also low-mass MS stars but the fact that no reflection effect in their light curves has been detected excludes this possibility (Maxted, Morales-Rueda & Marsh 2004; Kupfer et al. 2015). As we mentioned in the first paragraph of this section, it is possible that some WDs with masses smaller than 0.48 $\rm M_{\odot }$ in the sample are non-pure He WDs. For these kind of systems, the progenitor of WDs have masses larger than 2.0 M⊙ and non-degenerate He core ignition on the RGB. After their H envelopes are stripped, they can evolve into WDs with CO cores. With a binary population synthesis method, we estimate that the contribution of low mass (0.48 M⊙) non-pure He WD + sdB binaries to low mass (0.48 M⊙) WD + sdB binaries is less than $5{{\ \rm per\ cent}}$.
Citation: Han et al. 2002
Citation function: Uses; Uses