adsabs/FOCAL huggingface dataset:
To search for transit events, we will make use of our custom pipeline SHERLOCK (Pozuelos et al. 2020)5. This pipeline provides the user with easy access to Kepler, K2, and TESS data for both SC and LC. The pipeline searches for and downloads the pre-search data conditioning simple aperture (PDC-SAP) flux data from the NASA Mikulski Archive for Space Telescope (MAST). Then, it uses a multi-detrend approach in the WOTAN package (Hippke et al. 2019), whereby the nominal PDC-SAP flux light curve is detrended several times using a biweight filter or a Gaussian process, by varying the window size or the kernel size, respectively. This multi-detrend approach is motivated by the associated risk of removing transit signals, in particular, short and shallow signals. Each of the new detrended light curves, jointly with the nominal PDC-SAP flux, is then processed through the transit least squares (TLS) package (Hippke & Heller 2019) in the search for transits. In contrast to the classical box least-squares (BLS) algorithm (Kovács et al. 2002), the TLS algorithm uses an analytical transit model that takes the stellar parameters into account. Then, it phase folds the light curves over a range of trial periods (P), transit epochs (T0), and transit durations (d). It then computes the χ2 between the model and the observed values, searching for the minimum χ2 value in the 3D parameter space (P, T0, and d). The TLS algorithm has been found to be more reliable than the classical BLS in finding any type of transiting planet, and it is particularly well suited for the detection of small planets in long time series, such as those coming from Kepler, K2, and TESS. The TLS algorithm also allows the user to easily fine-tune the parameters to optimize the search in each case, which is particularly interesting for shallow transits. In addition, SHERLOCK incorporates a vetting module that combines the TPFplotter (Aller et al. 2020), LATTE (Eisner et al. 2020), and TRICERATOPS (Giacalone et al. 2021) packages, which allows the user to explore any contamination source in the photometric aperture used, momentum dumps, background flux variations, x–y centroid positions, aperture size dependences, flux in-and-out transits, each individual pixel of the target pixel file, and to estimate the probabilities for different astrophysical scenarios such as transiting planet, eclipsing binary, and eclipsing binary with twice the orbital period. Collectively, these analyses help the user estimate the reliability of a given detection.
Citation: Hippke & Heller 2019
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It has been claimed that, for solar-type stars, CA models could face some difficulties (in principle) explaining the presence of giant planets around metal-poor stars, or massive planets at long radial distances (see e.g., Helled et al. 2014). In particular, HR 8799 is a metal-poor A5 star included in our sample, hosting three giant planets orbiting beyond 10 AU (Marois et al. 2008). Then, some authors proposed that the planets orbiting HR 8799 are likely formed by GI, given that GI models can take place at large radii and in low metal environments (Marois et al. 2008; Dodson-Robinson et al. 2009; Meru & Bate 2010). Other works cast some doubts about the planet formation around HR 8799, showing that CA models are also possible (Currie et al. 2011) or even a combination of GI and CA (Marois et al. 2010). For the case of β Pictoris and HD 169142, also having giant planets at long distances, CA models seem to be favored (GRAVITY Collaboration 2020; Nowak et al. 2020; Pérez et al. 2019). However, an important word of caution about these works is in order. Some of the stars mentioned (HR 8799 and HD 169142) display (superficial) metal-poor abundances, showing in fact a λ Boötis pattern (see for example Fig. 8). As previously mentioned, the most accepted idea about the origin of this peculiar signature, suppose a solar-like composition for the original molecular cloud where the stars born, and then some kind of selective accretion to obtain a λ Boötis pattern. In this way, it would be not entirely appropriate to assume a metal-poor natal environment for stars like HR 8799, as assumed by some works to support a GI planet formation (e.g., Meru & Bate 2010). This fact was early noted by Paunzen et al. (2014), in their comparison of λ Boötis stars and Population II type stars. Numerical simulations of planet formation around λ Boötis stars should assume a solar-like composition (rather than a metal-poor natal environment), and this could have important consequences for the subsequent results.
Citation: Pérez et al. 2019
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The flux calibration was performed from two calibration sources: (1) the spectrum of the spectrophotometric standard star GD71 (Bohlin 2003), obtained in 2016 January, which is used to eliminate as much as possible any strong wavelength dependence; (2) the median combination of a set of 10 images of the standard star HZ 4 (Bohlin, Dickinson & Calzetti 2001) obtained right after the end of the cube observation with photometric conditions similar to the observation conditions. The exact value of the interferometer’s modulation efficiency (ME), which acts essentially as an additional throughput loss, is the major source of uncertainty on the absolute flux calibration (Martin & Drissen 2016). Interferometric images of the laser source have been obtained before and after the observation of the target in order to measure the variation of ME at the calibration laser wavelength (543.5 nm) with respect to its nominal value (85 per cent). We have measured a loss of 11.7 per cent. The initial flux calibration of M 31 in the SN3 filter has been corrected for this loss (thus multiplied by a factor 1.13). But given the possible uncertainty on this estimate, we have double-checked it using Hubble Space Telescope (HST) narrow-band images of the target. The advantage of HST’s narrow-band filters is that they can easily be simulated by integrating the spectral cube over the filter’s well-known transmission curve. As SITELLE’s cube flux is expressed in erg cm−2 s−1 Å−1 and given the filter transmission curve F(σ), the integrated flux $\tilde{\phi }_{\rm F}$, expressed in erg cm−2 s−1 Å−1, is
(15)
\begin{equation}
\tilde{\phi }_{\rm F} = \frac{\int _{-\infty }^{+\infty }\phi (\sigma )F(\sigma ){\rm d}\sigma }{\int _{-\infty }^{+\infty }F(\sigma ){\rm d}\sigma }\, .
\end{equation}
If one converts the flux in terms of surface brightness, expressed in erg cm−2 s−1 Å−1 per HST pixel surface (SHSTCamera), i.e.
(16)
\begin{equation}
\tilde{B}_{\rm F} = \frac{\tilde{\phi }_F }{S_{\rm SITELLE}}S_{\rm HSTCamera},
\end{equation}
with SSITELLE = 0.322 arcsec2, the image $\tilde{B}_{\rm F}$, once properly aligned, can be directly compared with the HST frame. We have made this comparison for three different regions of the FOV (called WFC3, ACS1 and ACS2, shown in Fig. 9) and four different filters (F656N, F658N, F665N and F660N). Histograms of the flux ratio between the integrated frames and the HST frames are presented in Fig. 10 and the first moments of their distributions are listed in Table 3.
Citation: Martin & Drissen 2016
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Galaxies are surrounded by vast gaseous haloes which extend well beyond the hosts’ stellar components: Early observations of quasar sight lines attributed the presence of absorption at multiple intermittent redshifts to gaseous haloes of intervening galaxies (e.g. Bergeron 1986; Bergeron & Boissé 1991; Lanzetta et al. 1995; Tripp, Savage & Jenkins 2000; Chen, Lanzetta & Webb 2001). In the past decade, owing to the rise of large spectroscopic surveys of galaxies with well-determined physical properties (e.g. SDSS), all sky UV surveys (e.g. GALEX), and improved sensitivity of UV spectrographs (e.g. COS), studies of the gaseous haloes of galaxies could systematically connect gas absorption properties to galaxy properties in statistically meaningful samples (e.g. Cooksey et al. 2010; Prochaska et al. 2011; Tumlinson et al. 2013; Liang & Chen 2014; Lehner, Howk & Wakker 2015). The aforementioned gaseous haloes are commonly referred to as the circum-galactic medium (CGM) and are ubiquitous in galaxies regardless of mass or star formation activity: even sub-L* galaxies (Bordoloi et al. 2014), and passive galaxies host a CGM (Thom et al. 2012). The current model of the CGM suggests the presence of a clumpy multiphase medium which extends beyond the virial radius of the host galaxy, with a declining radial density profile, containing a substantial amount of gas and metals (e.g. Werk et al. 2013, 2014; Liang & Chen 2014; Lehner et al. 2014, 2015; Prochaska et al. 2017). Observational studies targeting the CGM of L* galaxies showed that the CGM gas content is comparable to the mass of the interstellar medium (ISM; e.g. Chen et al. 2010; Tumlinson et al. 2011; Werk et al. 2014; Prochaska et al. 2017) and correlates positively with ISM properties (Borthakur et al. 2015). Additionally, CGM observations infer a significant amount of metals (e.g. Werk et al. 2013; Peeples et al. 2014) where CGM metallicities can extend to supersolar metallicities (Prochaska et al. 2017). The clumpy multiphase CGM consists of a warm gas T ∼ 104 − 5 K (clumpy in nature) embedded within a hot diffuse T ∼ 106 K medium (e.g. Heitsch & Putman 2009; Armillotta et al. 2017; Bordoloi et al. 2017). The multiphase structure of the CGM is corroborated by the variety of observed ionic species which survive at a vast range of temperatures: While the warm gas hosts the low ionization species (e.g. H i, Si ii, Si iii, C ii, C iv), the hot medium is home for the most highly ionized species (e.g. O vi, O vii). Additionally, the spectral line profiles of absorbers in the CGM can be reproduced by invoking a patchy medium (e.g. Stern et al. 2016; Werk et al. 2016), i.e. multiple high density gas clouds contribute to the optical depth along the line of sight thus leaving their kinematic imprint on the absorption line profile. For a review of the CGM, see Putman, Peek & Joung (2012) and Tumlinson, Peeples & Werk (2017).
Citation: Werk et al. 2013
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Citation: Werk et al. 2013
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Citation: Tumlinson et al. 2013
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Studying the dynamical properties of rotating neutron stars is a domain which brings out various interesting features when one assumes a perfect fluid. It is known that the centrifugal force of a rotating star counters gravitational force and hence one can expect massive stars to be fast rotors, at least in the initial stages of the stellar evolution. As a result of rotation a star may experience damping due to unstable oscillations such as the r-modes. The r-modes are one of many pulsating modes that exist in neutron stars and are characterized by the Coriolis force acting as the restoring force (Andersson 1998). The r-modes are unstable to emission of gravitational radiation (GR) by the Chandrashekhar-Friedman-Schutz (CFS) mechanism (Chandrasekhar 1970; Friedman & Schutz 1978). It was shown in Andersson (1998) that the r-modes are unstable for all rotating perfect fluid stars irrespective of their frequency. Dissipative effects such as shear and bulk viscosities work towards suppressing GR driven instabilities and has been studied by various authors over the past few years (Lindblom, Owen & Morsink 1998; Jones 2001; Lindblom & Owen 2002; van Dalen & Dieperink 2004; Drago, Lavagno & Pagliara 2005; Nayyar & Owen 2006; Jaikumar, Rupak & Steiner 2008; Jha, Mishra & Sreekanth 2010; Ofengeim et al. 2019) under various considerations. If the GR time-scale is shorter than the damping time-scale due to such dissipative processes, then the r-mode will be unstable and a rapidly rotating neutron star could lose a significant fraction of its rotational energy through GR. At higher temperatures (T > 109 K), the dominant dissipation is due to bulk viscosity, which arises due to density and pressure perturbations, a consequence of the star being driven out of equilibrium by oscillations. The system tries to restore equilibrium through various internal processes. In the case of r-modes, since the typical frequencies are of the order of the rotational frequencies of the stars, the reactions that dominate are the weak processes. Within these weak processes, although the modified Urca processes involving leptons are important, it has been shown that non-leptonic processes involving hyperons contribute more significantly towards bulk viscosity at temperatures lower than a few times 109 K (Lindblom & Owen 2002). Our goal here is to investigate the same using a chiral model calibrated to reproduce the desired nuclear matter properties, in particular the density content of the nuclear symmetry energy at both low and high densities.
Citation: Lindblom & Owen 2002
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Citation: Lindblom & Owen 2002
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By using the RADEX2 dense cloud models developed by Aalto et al. (2015) to reproduce the HCN(3–2)/(1–0) line luminosity ratios in the outflow of Mrk 231, we can attempt to find a combination of XHCN, XCN, Tkin, and nH2 solutions that can also fit the CN/HCN and CN spin doublet line ratios (Table 2). We assume that the HCN and CN line emissions arise from the same dense cloud population, while the low-J CO line emission is due to a different, more diffuse phase of the outflow. We recall that in these models (see also Aalto et al. 2015), the dense clouds can be either self-gravitating virialised clouds, which implies that their internal velocity dispersion (Δvsg) is locked to their mass (Mvir) and size (R) through Δvsg = (GMvir/G)1/2, or unbound clouds, for which Δv ≫ Δvsg. We explored CN and HCN abundances in the range between 10−8 and 10−6. We find that depending on whether the clouds are self-gravitating or unbound, the models produce very different values for the absolute CN and HCN abundances, hence XCN and XHCN remain quantitatively unconstrained for the outflow with current data. However, all possible solutions that fit the observed line ratios consistently require XCN >  XHCN, with a CN abundance that is at least a factor of three higher than the HCN abundance. Gas densities for this outflow phase (traced by the CN and HCN emissions) are nH2 ∼ 105 − 106 cm−3, with temperatures not much higher than Tkin ∼ 20 K. Because CN is a well-known PDR tracer (see also Sect. 1), these results strongly suggest that the whole dense cloud population in outflow is affected by UV radiation. We should mention that high CN abundances may also be due to cosmic rays (e.g. see work done on the Galactic centre by Harada et al. 2015), which are known to permeate the outflow of Mrk 231, as inferred by González-Alfonso et al. (2018) based on the high OH+ abundance. However, it is not clear whether a cosmic-ray chemistry would also explain XCN >  XHCN.
Citation: Aalto et al. (2015)
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Citation: Aalto et al. 2015
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It was found that the complex dielectric function from Pecharromán et al. (1999) for the sample obtained by heating bayerite at 1273 K, assuming a spheroid with depolarization parameters of (0.35, 0.003), produced an opacity with 11, 20, 28, and 32 μm features, so this component was included in the models. However, with only this component, the observed 20 μm features in the residual spectra were found to be wider than those in the models. By adding the opacity of the sample obtained by heating boehmite at 1173 K, the width of the 20 μm feature could be matched. This was done using the complex dielectric function for the sample obtained by heating boehmite at 1173 K from Pecharromán et al. (1999), assuming a spheroid with depolarization parameters of (0.35, 0.035). The complex dielectric functions of the samples obtained by heating bayerite and boehmite to various temperatures (Pecharromán et al. 1999) were derived by modeling the reflectance spectra of pellets obtained by pressing powders of these materials under great pressure. This method required Pecharromán et al. (1999) to assume an effective medium theory, such that a pellet is a mixture of one of their samples with a matrix of air. Pecharromán et al. (1999) noted that heating bayerite at 500°C eliminates the XRD pattern of bayerite, and they note that at 700°C, the infrared reflectance spectrum of the boehmite sample no longer shows OH− stretching bands. This must mean that the samples obtained from heating bayerite at 1273 K and from heating boehmite at 1173 K are no longer bayerite or boehmite, respectively. XRD performed by Pecharromán et al. (1999) of the sample of bayerite prepared at 1273 K suggests only θ-alumina was present, and their infrared and NMR spectroscopy confirms this. XRD of their sample obtained from heating boehmite to 1173 K (Pecharromán et al. 1999) suggests δ-alumina to be present, though some amounts of θ-alumina and α-alumina are present, as they deduce from XRD and infrared and NMR spectroscopy.
Citation: Pecharromán et al. (1999)
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Citation: Pecharromán et al. 1999
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Citation: Pecharromán et al. 1999
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Citation: Pecharromán et al. (1999)
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Citation: Pecharromán et al. (1999)
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Citation: Pecharromán et al. (1999)
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Citation: Pecharromán et al. (1999)
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The temperature gradient of the 3-equation model is comparable to results of different TCM approaches (Xiong & Deng 2001; Li & Yang 2007) for the base of the solar convective envelope. Both Zhang & Li (2012; their Figs. 6 and 7) and Xiong & Deng (2001; their Fig. 8) find a temperature gradient that transitions gradually from the adiabatic to the radiative value. They also find a Deardorff layer with a degree of sub-adiabaticity at the formal Schwarzschild boundary comparable to our findings. From the convective flux as presented in Xiong (1986) one also would expect a similar temperature gradient in the overshooting zone. Furthermore, the shape of the model temperature gradient is also in qualitative agreement with the discussion in Viallet et al. (2015). They argue that under the physical conditions in convective cores, in regions of overshooting efficient chemical mixing and a gradually transitioning temperature gradient are expected. In the 3-equation non-local model, the extent of the nearly adiabatic overshooting zone is controlled by the shape of the negative convective flux in the overshooting zone. For smaller (more negative) values of the convective flux (i.e. more efficient buoyancy braking) the temperature gradient is expected to be closer to the adiabatic value, while for larger (less negative) values it will be closer to the radiative temperature gradient. In Eq. (1) the negative convective flux and the dissipation term act as sink terms in the overshooting zone. Hence, the behaviour of the dissipation term will impact also on the convective flux and in turn on the value of the temperature gradient in the overshooting zone. In computations with the 1-equation non-local version of the theory, the negative convective flux is the dominant sink term for the TKE and the actual dissipation term is negligible (Fig. 8 of Paper I). This leads to more negative values of the convective flux and thus to a mostly adiabatic temperature gradient in the overshooting zone. We discuss this in more detail in Sect. 5.1 (we also refer to Fig. 10).
Citation: Xiong (1986)
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Since its discovery more than five decades ago (Cheung et al. 1968), ammonia (NH3) has been a most valuable molecule for investigating the physical properties of molecular clouds (e.g., Ho & Townes 1983). While thermally excited transitions in the centimeter-wavelength inversion transitions of ammonia are regarded as a reliable thermometer of molecular clouds (e.g., Walmsley & Ungerechts 1983; Danby et al. 1988), ammonia masers have attracted attention since the first detection of maser action in the (J, K) = (3,3) metastable (J = K) line toward the massive star-forming region W33 (Wilson et al. 1982). Subsequent observations have led to the detection of new metastable ammonia masers, including 15NH3 (3,3) (Mauersberger et al. 1986), NH3 (1,1) (Gaume et al. 1996), NH3 (2,2) (Mills et al. 2018), NH3 (5,5) (Cesaroni et al. 1992), NH3 (6,6) (Beuther et al. 2007), NH3 (7,7), NH3 (9,9), and NH3 (12,12) (Henkel et al. 2013). These have led to the discovery of metastable maser lines in 22 different regions (Mauersberger et al. 1986, 1987; Wilson & Henkel 1988; Wilson et al. 1990; Pratap et al. 1991; Cesaroni et al. 1992; Wilson & Schilke 1993; Mangum & Wootten 1994; Kraemer & Jackson 1995; Zhang & Ho 1995; Zhang et al. 1999; Walsh et al. 2007; Hunter et al. 2008; Galván-Madrid et al. 2009; Brogan et al. 2011; Urquhart et al. 2011; Walsh et al. 2011; Wang et al. 2012; Henkel et al. 2013; Hoffman & Joyce 2014; McEwen et al. 2016; Mills et al. 2018; Hogge et al. 2019; Mei et al. 2020; Towner et al. 2021). Compared with the metastable ammonia masers, detected non-metastable (J > K) ammonia maser transitions are more numerous. The first highly excited non-metastable ammonia maser was detected by Madden et al. (1986) in the (J, K) = (9,6) and (6,3) lines. Thereafter, many other NH3 non-metastable inversion transition lines have been identified as masers, including the (5,3), (5,4), (6,1), (6,2), (6,4), (6,5), (7,3), (7,4), (7,5) (7,6), (8,3), (8,4), (8,5), (8,6), (9,3), (9,4), (9,5), (9,7), (9,8), (10,7), (10,8), (10,9), and (11,9) transitions (e.g., Mauersberger et al. 1987, 1988; Walsh et al. 2007; Henkel et al. 2013; Mei et al. 2020). Except for the NH3 (3,3) masersproposed to be associated with four supernova remnants (McEwen et al. 2016), almost all the other ammonia masers are detected in high-mass star-forming regions (HMSFRs). However, while many HMSFRs host water (H2O), hydroxyl (OH), or methanol (CH3OH) masers, ammonia masers are quite rare in these sources, and the role that the environment of a young high-mass star plays in their excitation remains unclear. Therefore, dedicated searches for ammonia masers in HMSFRs are indispensable in regard to their overall incidence and association with different environments, which can provide additional constraints on the pumping mechanism of ammonia masers.
Citation: Mills et al. 2018
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Citation: Mills et al. 2018
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