adsabs/FOCAL huggingface dataset:
As a comparison, we also use the publicly available mock catalogue JAdes extraGalactic Ultradeep Artificial Realizations (JAGUAR; Williams et al. 2018) to derive the relation between UV and [O iii] + H β luminosity of simulated z ∼ 8 galaxies. The JAGUAR mock catalogue has been produced by matching luminosity and stellar mass functions as well as the relation between the stellar mass and UV luminosity, mostly at z ≤ 4. The galaxy properties are then extrapolated up to z ∼ 15. The JAGUAR catalogue provides emission line fluxes and EWs for the main lines based on modelling with the beagle code (Chevallard & Charlot 2016; Chevallard et al. 2018a). We identify all galaxies from the fiducial JAGUAR mock in the redshift range 7.11  z  9.05 and we randomly select 1000 of them to match the absolute UV magnitude distribution of our sample, and then fit the UV-[O iii] + H β luminosity data. The result is shown in red in Fig. 6. Similarly to our sample, the z ∼ 8 galaxies from the JAGUAR catalogue exhibit a tight relation between UV and [O iii] + H β luminosity (Spearman rank correlation coefficient ρ = 0.73, standard deviation from null hypothesis σ > 40). However, the mock galaxies exhibit a significantly lower [O iii] + H β luminosity (∼0.5 dex) at a given LUV compared to the relation of our galaxies. The detailed reason for this discrepancy relative to the JAGUAR mock is unclear, but one possible reason is differences in the median physical properties. For instance, while the mock galaxies exhibit (3.6–4.5)$\mu$m colour similar to the ones from our sample at a given UV luminosity, the average F125W-3.6$\mu$m colour in JAGUAR is smaller by ∼0.3 mag compared to the observed F125W-3.6$\mu$m colour in our sample. This means that while (3.6–4.5)$\mu$m colour and EW([O iii] + H β) are on average similar between JAGUAR and our sample, the absolute [O iii] + H β line luminosity scales with the 3.6$\mu$m flux which is larger in our sample compared to the JAGUAR mock catalogue. Furthermore, JAGUAR models a small field comparatively to our data, therefore the overlap in UV luminosity is small.
Citation: Chevallard et al. 2018a
Citation function: Uses
The progenitors of SNe Ib/c have been thought to be Wolf-Rayet (W-R) stars with high initial masses (MZAMS ≳ 25 M⊙; Crowther 2007). Before core collapse, these stars usually have experienced severe mass loss through strong stellar winds or due to interaction with companion stars (van der Hucht 2006; Paxton et al. 2015). As the evolution of massive stars is usually dominated by binary evolution (Heger et al. 2003) and also depends largely on metallicity, rotation, and so on (Heger et al. 2003; Georgy et al. 2013, 2012), this makes the direct identification of their progenitors complicated (Smartt 2015). However, there are increasing studies suggesting that a lower-mass binary scenario is more favorable for most SNe Ib/c, considering the measured low ejecta masses (Eldridge et al. 2013; Lyman et al. 2016). In addition, the H/He envelopes of the progenitor stars are stripped by binary interaction. There are many detections of progenitor stars for SNe II. For example, most SNe IIP are found to originate from red supergiants (Smartt et al. 2009), while SNe IIL are typically from progenitors with somewhat warmer colors (see Smartt 2015, for a review), and SNe IIb are from those with higher effective temperatures such as yellow supergiants that have had their H/He envelopes partially stripped through binary interaction (e.g., SN 1993J; Podsiadlowski et al. 1993; Maund et al. 2004; Fox et al. 2014). Until recently, there has been only one report of the possible identification of a progenitor star for SNe Ib, namely, iPTF 13bvn, which was proposed to spatially coincide with a single W-R-like star identified on the pre-explosion Hubble Space Telescope (HST) images (Cao et al. 2013; Groh et al. 2013). But such an identification is still controversial (e.g., Bersten et al. 2014; Fremling et al. 2014; Eldridge et al. 2015; Eldridge & Maund 2016). Direct detection of progenitor stars is still elusive for SNe Ic, which prevents us from further testing the theoretical evolution of massive stars (Eldridge et al. 2013).
Citation: Eldridge et al. 2013
Citation function: Background
Citation: Eldridge et al. 2013
Citation function: Background
We have presented MoCA, a Monte Carlo code for Comptonisation in Astrophysics which includes polarisation. To our knowledge MoCA is the first code operating with single photons and including all special relativity and quantum effects. The main disadvantage of this approach is the long computing time, which implies the need to parallelise the code on clusters of computers. The advantage with respect to pure analytical models such as those available in XSPEC is that we can explore the totality of the parameters space for the Comptonising medium (i.e. thermal energy and optical thickness of the corona) without any restriction and this approach will also allow a better understanding of the whole process. We also included all corrections such as Klein–Nishina cross-section and scattering angle distribution. These effects, small below 100 keV, must nonetheless be taken into account when inferring the thermal energy of the corona from observations. In some sources it has been found that coronae can have extremely high energy cut-off (e.g. NGC 5506, Matt et al. 2015) and therefore thermal energy, which is inferred by measuring the curvature of NuSTAR spectra at high energy and in this context K–N effects cannot be neglected. From the polarimetric point of view we did not seen any deviation due to K–N effects, but this was to expected as we focussed our attention below 100 keV where these effects are small. However, one can imagine a scenario in which the thermal energy of the corona is few tens of keV and in that case we expect to see a difference both on the spectrum and the polarisation but we defer such investigation to future papers focussed on the exploration of the coronal parameters space. In its actual form the code is fast enough to explore different geometries of the corona with different parameters. Spectra can then be compared with those obtained by NuSTAR to derive coronal parameters, especially in the high optical depth regime where analytical models are not reliable. As already mentioned, much observational evidence points in the direction of compact coronae above or around the compact object. In order to properly treat such coronae, gravitational effects must be taken into account. We have recently included a ray-tracing routine to take into account GR effects: this new version of MoCA, and applications of the code to different astrophysical scenarios, will be discussed in future papers. Nonetheless we have shown the geometrical effect of more compact coronae on the spectra and the polarisation signal: the spectra become softer as the corona shrinks and the polarisation changes dramatically as we approach a more symmetrical shape of the corona. The study we performed will also be useful to quantify the impact of GR effects on compact coronae with respect to a purely geometrical effect.
Citation: Matt et al. 2015
Citation function: Background; Background
ESA's Rosetta spacecraft orbited closely around comet 67P/Churyumov–Gerasimenko (hereafter 67P) during its 2 yr mission. From 2014 August to 2016 September, the camera system Optical, Spectroscopic, and Infrared Remote Imaging System (OSIRIS; Keller et al. 2007) onboard Rosetta observed the comet nucleus down to few centimetres per pixel providing detailed images of the comet's surface and activity. The study of 67P morphology has been performed with unprecedented detail thanks to the OSIRIS image spatial coverage and resolution. A first OSIRIS analysis on the nucleus structure, morphology and jets activity is reported in Sierks et al. (2015) and Thomas et al. (2015). The comet's surface revealed a wide diversity in morphology such as layering (Massironi et al. 2015), pits (Vincent et al. 2015), boulders and fractures (Pajola et al. 2015, 2016c; El-Maarry et al. 2015a), and high reflectivity boulders (Pommerol et al. 2015). During the entire mission, surface changes have been detected thanks to images acquired before and after perihelion passage (Groussin et al. 2015; Pajola 2017; El-Maarry 2017). The global spectrophotometric properties of the comet were investigated in detail by Fornasier et al. (2015) identifying three different groups of terrains depending on spectral slope values. In addition to a global characterization of the comet, OSIRIS allowed a detailed analysis of specific regions identified on the comet by means of high resolution and multifilter images (La Forgia et al. 2015; Deshapriya et al. 2016; Lucchetti et al. 2016; Oklay et al. 2016a; Pajola et al. 2016d; Oklay et al. 2017). We decided to follow the approach of characterizing at the highest resolution possible specific regions of interest (ROIs) on 67P. We therefore performed a detailed analysis on an area belonging to the Seth region (El-Maarry et al. 2015b) that is characterized by flat-floored and steep-walled circular depressions (El-Maarry et al. 2015b; Giacomini et al. 2016). We focused our attention on the circular niches of the Seth area (Fig. 1) and performed a multidisciplinary study of these structures investigating their geomorphological and spectrophotometric properties. In addition, thanks to images acquired pre- and post-perihelion, we conducted a comparative analysis to find if this area has been subjected to surface changes. This can be useful to provide constraints about Seth's niches properties as well as their possible origin being one of the interesting features located on the surface of 67P. Specifically, the circular niches deposits can be considered the result of landslide events that occurred on the comet surface, as the recent Aswan cliff collapse (Pajola 2017), where it has been reported the occurring of falling material from the adjacent cliff after the Rosetta perihelion passage. Hence, with this work we plan to understand if the formation of these Seth's circular niches occurred recently or if it is correlated to older events that have shaped the comet's surface.
Citation: Massironi et al. 2015
Citation function: Background
The CO emission, tracking the bulk of the gas, reveals an almost face-on one-armed spiral, of which almost two full windings can be traced. What could be the origin of this spiral structure? As the majority of AGB stars are in binary systems, and perhaps all host planets, interaction between the outflow and a sufficiently massive and nearby companion may be the explanation of the observed CO morphology. The intricate emission features in the inner 2″ of the central CO emission maps is strongly reminiscent of hydrodynamical simulations of wind–binary interaction by Mastrodemos & Morris (1998) and Mohamed & Podsiadlowski (2012), where the latter authors performed tailored simulations for the Mira AB system in which the outflow of the AGB star Mira A is perturbed by the presence of its close companion Mira B. The wind–binary interaction that ensues leads to what is known as wind Roche-lobe overflow (WRLOF), where the slow AGB wind is confined to the star’s Roche lobe, while overflowing through the L1 Lagrange point. Gravitational interaction of the overflowing material with the companion produces an intricate feedback system where the stellar outflow material is ejected into the surrounding CSE through two distinct streams (through L2 and the stagnation point3) which combine to form an annular stream. As this stream travels outwards, it creates the larger scale spiral observed in the wind. The morphology resulting from this particular type of wind–binary interaction is shown in Fig. 3 in Mohamed & Podsiadlowski (2012). In Fig. 11 we show the emission pattern seen in the central regions of the CO channel at υ*. We compare this image with the bottom left panel of Fig. 3 in Mohamed & Podsiadlowski (2012), an opacity map of the interaction zone. Though the two properties that are compared differ in nature, they likely still trace the same global morphological structure. Indeed, several of the predicted morphological features can be identified in the data of EP Aqr. The bright central region with a north and southward hook-like extension are strikingly similar, as are the eastern and western crescent-shaped “voids”, the overall shape, and the morphological properties of the small-scale instabilities.
Citation: Mohamed & Podsiadlowski (2012)
Citation function: Background
Detecting the halo response to the LMC-induced DM wake would be an exciting advancement in testing our assumptions about the properties of DM, as well as providing key constraints on the potential of the MW and the mass and orbital history of the LMC. However, the GC19 simulations give predictions for the response in the context of smooth MW DM and stellar halos. In reality, the MW stellar halo contains a wealth of substructure that is not yet phase mixed, in the form of stellar streams (e.g., Odenkirchen et al. 2001; Newberg et al. 2002; Belokurov et al. 2006; Grillmair 2006; Shipp et al. 2018; also see Newberg & Carlin 2016 for a recent review) and stellar clouds (e.g., Newberg et al. 2002; Rocha-Pinto et al. 2004; Jurić et al. 2008; Li et al. 2016). In addition, using a sample of MW halo main-sequence turnoff stars from the HALO7D survey (Cunningham et al. 2019a), Cunningham et al. (2019b) observed that the estimated parameters of the velocity ellipsoid (i.e., 





) were different in the different survey fields; these differing estimates could be interpreted as evidence that the halo is not phase mixed over the survey range (





 kpc). They also showed maps of the halo velocity anisotropy β in two halos from the Latte suite of FIRE-2 simulations (introduced in Wetzel et al. 2016), finding that the anisotropy can vary over the range 





 across the sky. Some of the variation in the β estimates appeared to correlate with stellar overdensities in the halos, indicating that galactic substructure is at least in part responsible for the different velocity distributions. While some substructure in the halo can be clearly identified as overdensities in phase space and removed from analysis, the presence of velocity substructure in the halo could complicate attempts to detect signatures of the LMC-induced DM wake. For example, Belokurov et al. (2019) recently argued that the Pisces Overdensity (Sesar et al. 2007; Watkins et al. 2009; Nie et al. 2015) might be stars in the wake trailing the LMC in its orbit, because of their net negative radial velocities. However, it remains difficult to conclusively argue this scenario given that these stars could also be in Galactic substructure (or, perhaps, stars that are in substructure and have been perturbed by the DM wake).
Citation: Shipp et al. 2018
Citation function: Background
LISA Pathfinder (LPF; Antonucci et al. 2011), a European Space Agency (ESA) mission that operated near the first Sun–Earth Lagrange point (L1) from 2016 January through 2017 July, is in an ideal orbit to make such measurements. However, LPF flew no instrumentation dedicated to micrometeoroid or dust detection. LPF’s primary objective was to demonstrate technologies for a future space-based observatory of millihertz-band gravitational waves. The key achievement of LPF was placing two gold-platinum cubes known as “test masses” into a freefall so pure that it was characterized by accelerations at the femto-g level (e.g., Armano et al. 2016, 2018b), the level required to detect the minute disturbances caused by passing gravitational waves. In order to reach this level of performance, the test masses were released into cavities inside the spacecraft and a control system was employed to keep the spacecraft centered on the test masses. This control system was designed to counteract disturbances on the spacecraft, including those caused by impacts from micrometeoroids. Shortly before LPF’s launch, it was realized that data from the control system, if properly calibrated, could be used to detect and characterize these impacts and infer information about the impacting particles (e.g., Thorpe et al. 2016). While such impact events have been reported by other spacecraft, LPF’s unique instrumentation makes it sensitive to much smaller and much more numerous impacts and allows the impact geometry to be more fully constrained. Early results from the first few months of LPF operations suggested that such events could indeed be identified and were roughly consistent with the pre-launch predictions of their effect on the control system (e.g., Thorpe et al. 2017). In this paper we present results from the first systematic search for micrometeoroid impacts in the LPF data set. Our data set consists of 4348 hr of data in both the nominal LPF configuration and the “Disturbance Reduction System” (DRS) configuration, in which a NASA-supplied controller and thruster system took over control of the spacecraft (Anderson et al. 2018). Our data set corresponds to the times when LPF was operating in a “quiet” mode, without any intentional signal injections or other disturbances. During this period, we have identified 54 impact candidates using our detection pipeline and manual vetoing. We have characterized the properties of this data set and compared it to several theoretical models for the underlying dust population.
Citation: Armano et al. 2016
Citation function: Background; Background
Close to the end of their formation, stars are surrounded by a gas and dust disk, from which planets form. Magnetic fields are known to play a key role in the star-disk system (Johns-Krull 2014). It is believed that the inner regions of the disk are significantly ionized by the stellar radiation and that accreting material flows along magnetic channels that connect the disk to the star (Koenigl 1991). Very long and intense X-ray flares in star-forming regions might occur in such long channels (Favata et al. 2005), but this is highly debated (Getman et al. 2008). These flux tubes might resemble those observed in the solar corona and diagnosed in the stellar coronae, but on a much larger scale. On the Sun we see the so-called coronal loops on the scale of several thousand kilometers in active regions, but some faint large-scale structures can extend up to ∼1 R⊙ (Reale 2014). Most solar flares occur in active region loops, but the long-lasting ones can involve more and more loops aligned in arcades. The other stars are so distant that we cannot resolve the flaring regions, but it is supposed that they occur in loops and even in arcades. Whereas the duration of solar flares typically ranges from a few minutes to several hours, stellar flares can be very intense, more than the solar bolometric luminosity, and long-lasting, including longer than one day, in very active stars. Several such gigantic coronal flares have been surveyed in star-forming regions (Favata et al. 2005) and where they occur is a big question. Magnetic instabilities in flux tubes were proposed to be the origin of the flaring activity also in T Tauri stars (Birk 1998; Birk et al. 2000), and long-lasting stellar flares might be expected to involve loop arcades (Getman et al. 2008), like those on the Sun. In long-lasting solar flares, the duration is mainly due to the progressive involvement of more and similar loops, therefore duration is not directly linked to the size of the flaring structures. This might also be the case for giant stellar flares. However, if a single stellar loop were flaring, the cooling time of the confined plasma would be proportional to the loop length (Serio et al. 1991; Reale 2014), and day-long flares would correspond to giant loops, as long as they possibly connected the star with the disk (Hartmann et al. 2016). There are ways to distinguish between a pure cooling in a single loop and a decay only due to progressive reduction of the energy release in a loop arcade (Reale et al. 1997), but the explanations are contested and the uncertainties are large (Getman et al. 2008). Several studies (Favata et al. 2005; Giardino et al. 2007) find results compatible with long magnetic channels in pre-main sequence (PMS) stars, but the derivation of the loop length is based on the assumption of a flare occurring in a single loop (Reale 2007).
Citation: Reale 2014
Citation function: Uses
The distance between the peak of the main source and the transient source (blue POSS I) is ≈6 pixels or ≈6 arcsec (≈0.3 kpc) towards the North, at an angle of ≈− 10° (Fig. 2; left-hand column; rows 4–5). From the PSF characteristics of the blue POSS I and blue subtracted POSS images (Section 2.1.1), the transient source is slightly resolved ( 2 × PSF FWHM). Because of the faintness of the transient and its proximity to the larger, brighter main source, an automated detection algorithm such as SExtractor (Bertin & Arnouts 1996), implemented in GAIA36 2016A, was unsuccessful in deblending the two sources. Therefore, a more rudimentary method was employed to determine the magnitude of the transient: using ds9, the transient and the brightest North-western source (Sections 2.2 and 2.3; Fig. 2; left-hand column; rows 4–5; Fig. 3) were fitted with ellipses to perform relative aperture photometry on the photometrically calibrated unconvolved POSS images, after sky subtraction (see below). This procedure shows that the transient source corresponds to ≳96 per cent of the flux of the North-western source (m ≈ 20.9 mag; USNO B1.0; Monet et al. 2003), which is equivalent to the transient being Δm ≲ 0.05 mag fainter than the North-western source (Fig. 2; left-hand column; rows 4–5). Similarly, the flux ratio between the main source and the transient is approximately a factor of 6 (Section 2.1.4). Given this very simplistic procedure to obtain the transient photometry, the magnitude of the transient is quite uncertain. A magnitude of m ≈ 21.0±0.2 mag will be adopted in the following discussion (Section 3), with the confidence interval provided assuming a 20 per cent error in the flux (detection at a 5σ level; see below), corresponding to the estimated photometric error at the maximum intensity extended to the full source, and neglecting unquantifiable systematics errors; the error estimation is included in AppendixA. Favourably, changes in the magnitude as large as Δm ≈ 1.5 will not significantly alter the conclusions of this work (Section 3).
Citation: Bertin & Arnouts 1996
Citation function: Uses; Compare/Contrast