adsabs/FOCAL huggingface dataset:
To overcome these limitations, in our previous paper (Chatzistergos et al. 2018b, Paper I, hereafter) we introduced a novel approach to process the historical and modern Ca II K observations, to perform their photometric calibration, to compensate for the intensity centre-to-limb variation (CLV, hereafter), and to account for various artefacts. By using synthetic data, we also showed that our method can perform the photometric calibration and account for image artefacts with higher accuracy than other methods presented in the literature. More importantly, we showed that, as long as the archives are consistent with each other, for example, they are centred at the same wavelength and employing the same bandwidth for the observations, the method can be used to derive accurate information on the evolution of plage areas without the need of any adjustments in the processing of the various archives (Chatzistergos et al. 2019b, Paper II, hereafter). In Paper II, we applied our method to 85 972 images from 9 Ca II K archives to derive plage areas and produce the first composite of plage areas over the entire 20th century. In particular, we analysed the Ca II K archives from the Arcetri, Kodaikanal (8-bit digitisation), McMath-Hulbert, Meudon, Mitaka, Mt Wilson, Rome/PSPT, Schauinsland, and Wendelstein sites. Five out of the nine analysed archives were amongst the most studied and prominent ones, while the remaining archives were from less studied data sources. There are, however, many other Ca II K archives that are available and still remain largely unexplored. These archives harbour the potential to fill gaps in the available plage series as well as to address inconsistencies among the various archives and within individual archives (e.g. change in data quality, or in the measuring instrument with time). Moreover, since the work presented in Paper II, more data from various historical and modern archives became available in digital form. In particular, historical data that have been made available in the meantime include those from the latest 16-bit digitisation of the Kodaikanal archive, Catania, Coimbra, Kenwood, Kharkiv, Kyoto, Manila, Rome, Sacramento Peak, and Yerkes observatories, as well as additional data from the Meudon and Mt Wilson archives. In this light, here we present results from the most comprehensive analysis to date of historical and modern Ca II K observations taken between 1892 and 2019 from 43 different datasets for the purposes of producing a composite plage area series.
Citation: Paper II
Citation function: Extends; Extends
We made use of the Pipe3D pipeline (Sánchez et al. 2016a,b, 2018) to perform the spatially resolved stellar population analysis of the data cubes. We note that Pipe3D is based on the stellar population synthesis code FIT3D (Sánchez et al. 2016a,b), which uses the GRANADA and MILES simple stellar population (SSP) libraries in the current implementation to model and subtract the stellar spectrum and fit the emission lines. The GRANADA and MILES (gsd156) stellar library is a combination of the empirical stellar library of Vazdekis et al. (2010) and the theoretical stellar libraries of González Delgado et al. (2005) and González Delgado & Cid Fernandes (2010). The SSPs use a Salpeter (1955) initial mass function and cover 39 stellar ages (from 1 Myr to 14.1 Gyr1) and four metallicities (Z = 0.002, 0.008, 0.02, and 0.03). It is important to note that FIT3D performs the SSP fitting within the wavelength of 3500 Å –7000 Å. In addition, FIT3D considers the effects on dust extinction during the stellar population synthesis using a Cardelli et al. (1989) extinction law. FIT3D provides the SSP decomposition of the modeled stellar spectra of each spatially resolved region of the galaxies. With this decomposition, we can estimate the respective SFHs and the cumulative mass temporal distributions (or mass growth histories; e.g., Ibarra-Medel et al. 2016, 2019; Sánchez et al. 2019). In order to estimate the contemporary (local and integrated) SFR values, we smoothed the SFHs within a period of 30 Myr. These quantities are independent of SFR estimations that use the Hα emission line flux. For star-forming galaxies, both estimates are very similar, showing that the SFR calculated in the last 30 Myr is related to massive young stars, contributing to the Hα emission. For retired galaxies, the SFR calculated from the Hα flux is likely overestimated because this flux is dominated by effects of post-AGB stars rather than of young stars (see, e.g., Bitsakis et al. 2019; Sánchez et al. 2019). However, in these galaxies, the SFR derived from the stellar population synthesis with Pipe3D should also be considered as an upper limit. To fit the emission lines, Pipe3D subtracts the modeled stellar spectra and fits a series of Gaussian profiles for each line (Sánchez et al. 2016b). With those fits, Pipe3D provides the total flux, dispersion, and the kinematics of each emission line. The quality of the emission line fits of Pipe3D (and other codes) was tested using MUSE data (Sánchez-Menguiano et al. 2018; Bellocchi et al. 2019; López-Cobá et al. 2020). With the use of the adjacent continuum windows (from the modeled stellar spectra) for each emission line, Pipe3D estimates the stellar continuum flux and provides the equivalent width of each emission line.
Citation: Ibarra-Medel et al. 2016
Citation function: Uses
The observed QPOs are generally attributed to oscillations of the star excited by an explosion of magnetic origin that creates the flare. The oscillating stellar surface should modulate the charge density in the magnetosphere, creating variations in the optical depth for resonant Compton scattering of the hard X-rays that accompany the flare (Timokhin et al. 2008; D’Angelo & Watts 2012). In this connection, the problem of finding the oscillatory modes for a strongly magnetized neutron star has received much attention, and has proven to be a formidable problem. To make the problem tractable, most theoretical treatments of the QPO problem have assumed smooth field geometries, usually dipolar or variants (e.g., Glampedakis et al. 2006; Levin 2006, 2007; Sotani et al. 2008a, 2008b; Cerdá-Durán et al. 2009; Colaiuda et al. 2009; Colaiuda & Kokkotas 2011; Gabler et al. 2011, 2012, 2013a, 2013b, 2014; van Hoven & Levin 2011, 2012; Passamonti & Lander 2013). Smooth field geometries support a problematic Alfvén continuum that couples to the discrete natural spectrum of the crust. As pointed out by Levin (2006), if energy is deposited in the crust at one of the natural frequencies of the crust, and this frequency lies within a portion of the core continuum, the energy is lost to the core continuum in less than 0.1 s as the entire core continuum is excited. The crust excitation is effectively damped through resonant absorption, a familiar process in MHD; see, e.g., Goedbleod & Poedts (2004). The problem has been addressed by assuming field geometries with gaps in the Alfvén continuum. Under this assumption, long-lived quasi-normal modes can exist inside the gaps or near the edges of the Alfvén continuum. van Hoven & Levin (2011) showed for a “box” neutron star that introduction of a magnetic tangle breaks the Alfvén continuum. Link & van Eysden (2015) showed that for magnetic tangling in a spherical neutron star the problematic Alfvén continuum disappears.3

3
Sotani (2015) added general relativity in the treatment of the magnetic field, and confirmed some of the results of Link & van Eysden (2015).
 They found that the star acquires discrete normal modes, and quantified the mode spacing. It is clear from these investigations that the unknown magnetic field geometry is the most important ingredient in determining the oscillation spectrum of a magnetar.
Citation: Gabler et al.; 2013a
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To understand the 3D magnetic properties of the FC observed here, we reconstructed the coronal magnetic field of the source region based on the nonlinear force-free field (NLFFF) assumption using the HMI photospheric vector magnetogram as the bottom boundary (Yan et al. 2001; Canou et al. 2009; Wiegelmann & Sakurai 2012; Yang et al. 2016b; Zhong et al. 2019; Qiu et al. 2020). We first tried the extrapolation method (Guo et al. 2010; Zhu et al. 2016) and found that it is difficult to reproduce the magnetic field comparable with the morphology of the observed filament threads. We then resorted to the magnetic flux rope (MFR) embedding method based on regularized Biot-Savart laws (RBSL; Titov et al. 2014, 2018; Guo et al. 2019). The procedure was divided into four steps. First, we took advantage of the 304 Å images at 05:24 UT to derive the path of the FC, as shown in Fig. 7a. According to previous statistics (Tandberg-Hanssen 1995; Filippov & Den 2000; Engvold 2015), the height of active region filaments ranges from 5 Mm to 30 Mm. We therefore set the height of the MFR axis to 30 Mm and the MFR minor radius to 25 Mm, assuming that the lower half of the MFR is fully filled by cool material. The 3D path of the MFR was estimated according to Guo et al. (2021a). Second, we calculated a potential field utilizing the normal magnetic field component, where the projection effect was corrected (Wiegelmann et al. 2006; Guo et al. 2017). Third, we set the physical parameters of the RBSL model including the average value of unsigned magnetic flux at the two MFR footprints (1.72 × 1021 Mx) and the strength of the electric current following Eq. (12) in Titov et al. (2018). Finally, we inserted the MFR derived by the RBSL model into the potential field along the path of the FC and then performed a relaxation using the magnetofrictional code (Guo et al. 2016a,b). After relaxation, the force-free metric was σJ = 0.28, and the divergence-free metric was ⟨|fi|⟩ = 1.47 × 10−4, which were sufficiently small and generally acceptable according to Guo et al. (2021a)1.
Citation: Guo et al. 2019
Citation function: Uses
The theory of angular momentum evolution due to mass transfer via Roche lobe overflow and gravitational wave radiation is well established and leads to quantitative predictions (Landau & Lifschitz 1958; Paczyński 1967; Faulkner 1971; Verbunt 1993). However, angular momentum losses due to emission of gravitational waves only dominate at short orbital period (Porb ≲ 2 h), while other mechanisms, such as magnetic braking, are invoked to explain the orbit evolution at longer periods. The theory at the heart of the magnetic braking model is based on the same physical principles through which magnetic stellar winds are known to decelerate the rotation of low-mass stars (e.g. Kraft 1967; Skumanich 1972; Soderblom 1983). In synchronized binaries (typically the case of LMXBs), the angular momentum loss from the companion star is continuously re-distributed to the angular momentum of the whole binary system, therefore affecting its orbital period and binary separation (e.g. Tauris 2001). Large uncertainties are involved in the extrapolation of magnetic braking from isolated low-mass stars to the case of synchronized companion stars in LMXBs (e.g. see discussion in Knigge, Baraffe & Patterson 2011). Indeed, despite the large allowed range of model parameters (possibly translating into more than one order of magnitude on efficiency of angular momentum removal), magnetic braking often fails to reproduce measured values of the orbital period derivative, at least in the case of conservative mass transfer (Di Salvo et al. 2008; Hartman et al. 2008, 2009; Hu, Chou & Chung 2008; Burderi et al. 2010; Jain, Paul & Dutta 2010; González Hernández, Rebolo & Casares 2012, 2014). Additional mechanisms, typically involving outflows, are generally necessary to explain these discrepancies (Di Salvo et al. 2008; Hartman et al. 2008, 2009; Burderi et al. 2010). We also note that some of the best-monitored sources (i.e. EXO 0748−676; Wolff et al. 2009) show erratic variations of the orbital period with amplitudes of ΔP/P ∼ 10−5 and possible trends recurring on time-scales of years-to-decades, such as those observed in e.g. Algols systems (Kreiner & Ziolkowski 1978; Hall 1989). These variations are thought to be the consequence of the torque provided by the magnetic activity of a sub-surface magnetic field in the companion star with convective envelope. The latter induces a cyclic exchange of angular momentum between the inner and outer parts of the companion star causing a change in the gravitational quadrupole moment (Applegate 1992; Applegate & Shaham 1994; Lazaridis et al. 2011).
Citation: Tauris 2001
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The origin of the observed colour bimodality must be explained by successful models of galaxy formation. Considerable efforts have been directed towards reproducing the observed distribution of galaxy colours using different semi-analytic models of galaxy formation (Menci et al. 2005; Cattaneo et al. 2006, 2007; Driver et al. 2006; Cameron et al. 2009; Trayford et al. 2016; Nelson et al. 2018; Correa, Schaye & Trayford 2019). A significant number of studies have also been devoted to understand the spatial distribution of red and blue galaxies using various clustering measures like two-point correlation function (Zehavi et al. 2011), three-point correlation function (Kayo et al. 2004), genus (Hoyle et al. 2002), filamentarity (Pandey & Bharadwaj 2006), and local dimension (Pandey & Sarkar 2020). The studies of mass function of red and blue galaxies (Drory et al. 2009; Taylor et al. 2015) play a crucial role in guiding the theories of galaxy formation and evolution. Any such analysis requires one to classify the red and blue galaxies using some operational definition. But colour is a subjective judgment and it is difficult to classify galaxies according to their colours in an objective manner. Also, one should remember that it is the observed bimodal distribution of galaxy colours that motivates us to classify the galaxies according to their colours. In reality, no galaxies can be regarded as truly either ‘red’ or ‘blue’ based on their colours. Strateva et al. (2001) prescribed that the red and blue galaxies can be separated by using an optimal colour separator (u − r) = 2.22. However, a substantial overlap between the two populations prohibit us to objectively define galaxies as ‘red’ and ‘blue’ based on their colours. The boundary separating the two populations remains arbitrary and currently there are no meaningful theoretical arguments favouring any specific cut over others. Applying a hard cut to define samples of red and blue galaxies is expected to introduce significant contamination in these samples. Previously, Coppa et al. (2011) used an unsupervised fuzzy partition clustering algorithm applied on the principal components of a PCA analysis to study the bimodality in galaxy colours. Another study by Norman et al. (2004) use Bayesian statistical methods instead of sharp cuts in parameter space for such classification problems.
Citation: Zehavi et al. 2011
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Recently, dust-obscured galaxies (DOGs, Dey et al. 2008) shed light on this issue, because SMBHs in DOGs are expected to be rapidly growing during the coevolution. DOGs are originally defined as galaxies that are bright in mid-infrared (MIR), while faint in optical. Specifically, 





 (i.e., 





; Dey et al. 2008; Fiore et al. 2008). In the context of the gas-rich major-merger scenario, it is expected that the SF phase evolves into the AGN phase because the merging event leads to the active SF, while the gas accretion onto the nucleus caused by such a merger requires some time (see, e.g., Davies et al. 2007; Hopkins 2012; Matsuoka et al. 2017). Because such active galaxies are expected to be heavily surrounded by dust, DOGs potentially correspond to galaxies in the SF phase or AGN phase (Dey et al. 2008; Hopkins et al. 2008). DOGs are classified into two subclasses according to their spectral energy distribution (SED): “bump DOGs” and “power-law (PL) DOGs” (Dey et al. 2008). The bump DOGs show a rest-frame 1.6 μm stellar bump in their SEDs, while the PL DOGs show a power-law feature on their SEDs. Therefore, it is considered that the bump DOGs correspond to galaxies in SF mode (Desai et al. 2009; Bussmann et al. 2011), while the PL DOGs correspond to galaxies in the AGN phase (Fiore et al. 2008; Bussmann et al. 2009; Melbourne et al. 2012). The fraction of PL DOGs among all DOGs increases with increasing MIR flux density (e.g., Dey et al. 2008; Toba et al. 2015), which is similar to the behavior of the luminosity dependence of the AGN fraction in ultraluminous infrared (IR) galaxies (ULIRGs; see Sanders & Mirabel 1996 for a review). The comoving number density of DOGs shows its peak at z ∼ 1–2 (e.g., Dey et al. 2008; Toba et al. 2017), which corresponds to the peak of SF rate density and the growth rate of SMBHs (e.g., Richards et al. 2006; Madau & Dickinson 2014). This strongly suggests that DOGs are related to the most active objects in terms of the coevolution between galaxies and SMBHs. In this sense, DOGs with a high IR luminosity potentially harbor a rapidly growing SMBH and are, therefore, important for understanding the coevolution of galaxies and SMBHs.
Citation: Hopkins 2012
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However, using the well-assessed new post-AGB evolutionary models, we confined the main-sequence ages of our halo sample to be mostly ∼2–5 Gyr, with the oldest being ∼6–8 Gyr, while the outer-disk sample are mostly ≲1–4 Gyr. We thus conjecture that our targets probably formed prior to the encounter with M33. Obviously, our sample represents the population that is different from the underlying, smooth, extended (and mostly metal-poor) halo component of M31 (Ibata et al. 2007, 2014), which was formed through the repeated accretion of smaller galaxies in the distant past. These bright PNe seem to resemble the younger, metal-rich population in the outer stream of M31, as revealed by HST pencil-beam pointings on the Giant Stream (Brown et al. 2006a; Bernard et al. 2015b). The metallicity of the stream fields was enriched continuously from [Fe/H] ∼ −1.5 to at least solar level about 5 Gyr ago (Bernard et al. 2015b). This timeline of metal enrichment is generally consistent with the stellar ages of our metal-rich sample. N-body simulations suggested that the Giant Stream and other stream-like features in the halo are debris of a massive (≳109–





) progenitor that was recently disrupted during the course of a merger (e.g., Ibata et al. 2004; Fardal et al. 2006, 2007, 2008, 2013; Font et al. 2006; Geehan et al. 2006; Mori & Rich 2008; Sadoun et al. 2014). The extended star formation history and the broad range of metallicity (−1.5 ≲ [Fe/H] ≲ 0.2) discovered in the stream fields can be explained by a disk galaxy progenitor (Brown et al. 2006a, 2006b; Bernard et al. 2015b). If the stellar streams in M31's halo indeed have a common origin, our sample of halo PNe then probably formed through extended star formation in this possibly massive, disk-like progenitor. Moreover, some simulations predict that the remnant of the disrupted satellite resides in the NE Shelf (e.g., Fardal et al. 2008, 2013; Sadoun et al. 2014); PN17 in our sample is located in this region and might be associated with this substructure (see Section 4.4).
Citation: Bernard et al. 2015b
Citation function: Background; Similarities
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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Fast radio bursts (FRBs) are bright, cosmological origin, and millisecond-duration bursts in radio wavelengths (Lorimer et al. 2007; Thornton et al. 2013; Bassa et al. 2017; Macquart et al. 2020). After the discovery of the first FRB (Lorimer et al. 2007), a number of dedicated facilities have been conducted to search FRBs, such as the Parkes telescope (e.g., Bhandari et al. 2018), the updated Molonglo Observatory Synthesis Telescope (e.g., Farah et al. 2018), the Australian Square Kilometre Array Pathfinder (e.g., Shannon et al. 2018), the Canadian Hydrogen Intensity Mapping Experiment (CHIME; The CHIME/FRB Collaboration et al. 2018), the Deep Synoptic Array (Kocz et al. 2019; Ravi et al. 2019), the Green Bank Telescope (Masui et al. 2019), Arecibo (Spitler et al. 2014; Patel et al. 2018), and the Five-hundred-meter Aperture Spherical radio Telescope (FAST; Nan et al. 2011; Li et al. 2019). All these efforts result in an increasing rate of new FRB detections. Among them, more than 20 repeating FRBs have been reported. Particularly, the physical origin of the repeating FRB 20121102A was identified to be with a low-metallicity star-forming dwarf galaxy at a redshift 0.19273 (Bassa et al. 2017; Tendulkar et al. 2017). Another repeating FRB 20190523A was found to be associated with a more massive but low specific star formation rate (Ravi et al. 2019). The identification of the counterpart of the brightest radio bursts from SGR 1935+2154 as a magnetar in our Galaxy by HXMT (Li et al. 2021) and INTEGRAL (Mereghetti et al. 2020) with short-duration X-ray bursts suggests that at least a fraction of FRBs are connected with newborn magnetized neutron stars (e.g., Weltman & Walters 2020; Zhang 2021). More bursts with similar characteristics need to be detected in the future to confirm this conclusion. Recently, a new large sample with 535 FRBs was presented by CHIME/FRB Collaboration et al. (2021) that were detected by the CHIME survey, including 61 bursts from 18 previously reported repeating sources and 474 one-off bursts. Though an increasing catalog of theories and models is developing to explain the physical nature of FRBs (e.g., see the review of Platts et al. 2019; Xiao et al. 2021), the origin of FRBs remains a mystery.
Citation: Spitler et al. 2014
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