adsabs/FOCAL huggingface dataset:
There are additional lines of evidence suggesting that accretion in HD 100546 could be magnetospheric, and not through a BL. Equation 3 in Johns-Krull, Valenti & Koresko (1999) provides a lower limit to the magnetic field necessary to drive MA, in terms of the stellar parameters and accretion rate. Assuming a mass accretion rate of ∼10−7 M⊙ yr−1, the minimum stellar mass and maximum stellar radius allowed by the uncertainties provided by Fairlamb et al. (2015), and a minimum rotation period of 0.26 d (from a maximum projected rotational velocity vsin i = 65 km s−1 and a minimum inclination i = 22°, from Guimarães et al. 2006, and TAT11), HD 100546 could require a magnetic field of only several tens of Gauss to drive accretion magnetospherically.1 This is consistent with the magnetic field measured by Hubrig et al. (2009), 89 ± 26 G, although non detections have also been reported (Donati et al. 1997; Hubrig et al. 2013). In addition, if accretion in HD 100546 is actually magnetospheric the Keplerian gaseous disc should be truncated at a stelleocentric radius that increases with the stellar magnetic field and radius, and decreases with the accretion rate and stellar mass (Elsner & Lamb 1977). Using equation 6 in Tambovtseva et al. (2014), the disc truncation radius of HD 100546 should be ≲ 0.01 au (Fig. 3). This value is consistent with the Keplerian radius inferred from the width of the Br γ line profile (∼ ± 200 km s−1; Section 3); once this value is de-projected using the inclination in TAT11, the corresponding Keplerian distance is ≲ 0.02 au. It is noted that, in contrast to optically thick lines like H α, Br γ is mainly broadened by the Doppler effect (Tambovtseva et al. 2014). If the Keplerian disc is truncated by the stellar magnetic field, the gas would then fall ballistically on to the stellar surface. This could eventually be traced from the presence of redshifted absorptions at free-fall velocities in the profiles of several lines. In fact, previous spectroscopic analysis involving optical/UV lines suggest MA/ejection processes in HD 100546 (Vieira et al. 1999; Deleuil et al. 2004), with redshifted absorptions at velocities comparable to free-fall (Guimarães et al. 2006). However, those signatures are not observed in the Br γ profile of HD 100546. The unresolved component contributing ∼35 per cent to the visibility (Section 4.1) could be related to magnetospheric infall, but the small spatial scales involved cannot be probed from our observations.
Citation: Hubrig et al. 2013
Citation function: Similarities
The earliest observations of magnetic fields in astrophysics were sunspot field strength measurements at the Mount Wilson Observatory (MWO) in California, USA, in 1908 (Hale 1908). Observations employed the Zeeman effect and were based on measuring the separation (splitting) between the two components of a spectral line, first Fe I 6173 Å and later Fe II 5250 Å. Daily observations of sunspot magnetic fields have been conducted since 1917 (Hale et al. 1919). In the early 1950s, the invention of an electronic magnetograph (Babcock 1953) allowed the measurement of regions with weaker magnetic fields than in sunspots, such as plages. Regular full-disk magnetograms have been observed since early 1960, first at MWO (Howard 1974) and then at the National Solar Observatory (NSO) at Kitt Peak starting in 1973 (Livingston et al. 1976). In our long-term project aimed at reconstructing the past magnetic activity on the Sun (Pevtsov et al. 2016, 2019; Virtanen et al. 2017, 2018, 2019a,b), we also use indirect measurements of magnetic fields obtained, in particular, from chromospheric spectroheliograms. Recently, Chatzistergos et al. (2021) and Shin et al. (2020) have used chromospheric observations to reconstruct the solar magnetic activity in the past. Reconstructing past magnetic fields is important for understanding the long-term behavior of the Sun since it is the main factor affecting space climate. Ca II K spectroheliograms are essential for this task, since the first ones were taken already in the 1890s in Europe (Paris and Meudon Observatories, France; Deslandres 1909; Malherbe & Dalmasse 2019) and in the United States (Kenwood Observatory; Hale 1893). Continuous observation campaigns in the Ca II K line started in the early 20th century in Kodaikanal (India), Mount Wilson Observatory (USA), the National Observatory of Japan (Japan), Paris-Meudon Observatory (France), the Arcetri Astrophysical Observatory (Italy), and the Astronomical Observatory of Coimbra University (Portugal) (for a historical overview see, e.g., Bertello et al. 2016; Chatzistergos et al. 2020a). Chromospheric spectroheliograms, together with flux transport simulations, have also been successfully used to study the evolution of large-scale solar magnetic fields, particularly the polar fields, which are only partially observable (Virtanen et al. 2019a).
Citation: Virtanen et al. 2017
Citation function: Motivation; Motivation
Finally, the multiphase nature of galactic outflows implies that measurements of the outflow properties based on a single gas phase can lead to misleading conclusions (for a discussion, see e.g., Cicone et al. 2018b). Historically, systematic studies of galactic outflows in nearby and high-z galaxies have focused on the ionized gas – for example, as observed as broad wing emission in the spectra of the Hα, [O III] or Paα lines – (e.g., Heckman et al. 1990; Rupke & Veilleux 2013a; Woo et al. 2016; Harrison et al. 2016; Förster Schreiber et al. 2019; Ramos Almeida et al. 2019) and the atomic phase – based on the Na D or Mg II lines in absorption (e.g., Heckman et al. 2000; Rupke et al. 2002, 2005; Weiner et al. 2009; Roberts-Borsani & Saintonge 2019). The molecular component of outflows, on the other hand, has been much more difficult to study. Great progress was made with the Herschel Space Observatory using the OH 119 μm line in absorption to study molecular outflows in Seyfert and luminous infrared galaxies (Fischer et al. 2010; Sturm et al. 2011; Veilleux et al. 2013; Bolatto et al. 2013; Spoon et al. 2013; George et al. 2014; Stone et al. 2016; González-Alfonso et al. 2017; Zhang et al. 2018). More recently, the advent of powerful millimeter-wave interferometers such as the Atacama Large Millimeter/submillimeter Array (ALMA) and the NOrthern Extended Millimeter Array (NOEMA) are rapidly increasing the number of molecular outflows detected based on observations of the CO line (e.g., Combes et al. 2013; Sakamoto et al. 2014; García-Burillo et al. 2014; Leroy et al. 2015; Feruglio et al. 2015; Morganti et al. 2015; Dasyra et al. 2016; Pereira-Santaella et al. 2016, 2018; Veilleux et al. 2017; Fluetsch et al. 2019; Lutz et al. 2020). At high-z, so far only a handful of large-scale, molecular outflows have been studied in QSOs (e.g., Cicone et al. 2015; Vayner et al. 2017; Feruglio et al. 2017; Carniani et al. 2017; Fan et al. 2018; Brusa et al. 2018), sub-millimeter galaxies (e.g., Spilker et al. 2018), and main-sequence, star-forming galaxies (e.g., Herrera-Camus et al. 2019).
Citation: Rupke et al.; 2005
Citation function: Background; Background
Citation: Heckman et al. 2000
Citation function: Background; Background
Several studies have discussed the radio luminosity of GLEAM-X J1627 during its radio outburst in comparison with the limits of its rotational energy (Erkut 2022; Hurley-Walker et al. 2022). In particular, assuming isotropic emission, the radio luminosity of the brightest single peaks (L
radio ∼ 1030–1031 erg s−1; Hurley-Walker et al. 2022) exceeds the limits on the rotational power of the source by a few orders of magnitude. Figure 6 shows those peak radio luminosities and the rotational power of GLEAM-X J1627 in comparison with other pulsars, rotating radio transients (RRATs) and radio-loud magnetars. For the radio-loud magnetars, given their large variability, we have chosen the brightest radio pulses reported in the literature (data collected from Camilo et al. 2006, 2007; Weltevrede et al. 2011; Deller et al. 2012; Lynch et al. 2015; Majid et al. 2017; Pearlman et al. 2018; Lower et al. 2020, and Esposito et al. 2021). It is well known that assuming isotropic radio emission is not realistic, and a beaming factor necessarily has to be present (see, e.g., Erkut 2022). However, the relation between the duty cycle and the spin period of canonical pulsars has a large spread (Manchester et al. 2005). Moreover, it is observed that this relationship does not apply to radio-loud magnetars, which in general show larger duty cycles than what one would expect from the extrapolation of this tentative relation for radio pulsars to magnetars (see, e.g., Camilo et al. 2006, 2007). To avoid the uncertainty of beaming models, which for magnetars are mostly unknown even theoretically, we plotted the isotropic radio luminosity for all the different pulsar classes in Figure 6. From this plot, at variance with canonical radio pulsars, we see how the brightest single peaks for radio-loud magnetars might exceed their rotational powers, in line with what is possibly observed for GLEAM-X J1627. While not resolving uncertainties related to the exact mechanism of radio emission or the beaming factor, Figure 6 shows that, under the assumption of isotropic emission, even for magnetars the brightest single peaks exceed their rotational energy budget. Considering all the uncertainties in the assumptions used to derive the radio luminosities plotted in Figure 6, GLEAM-X J1627’s radio luminosity excess over its rotational power cannot be used as an argument for or against its neutron star nature.
Citation: Camilo et al. 2006
Citation function: Uses
Citation: Camilo et al. 2006
Citation function: Compare/Contrast
In the optical range, several proxies can serve as indicators of LyC leakage. In particular, line ratios involving ions with different ionization potentials, produced at different depths, have been proposed as indicators to discriminate between radiation-bounded and density-bounded H II regions. Radiation-bounded regions correspond to ionized spheres delimited by their Strömgren radii set by the equilibrium between production of photons by stars and ionization of the surrounding gas. Density-bounded regions are instead delimited by the lack of matter, which sets their outer radius before the Strömgren radius. Hence, they allow part of the LyC-photons produced by stars to escape from H II regions. The oxygen line ratio [O III]λ 5007 Å/[O II] λλ3726, 3728 Å (O32) proposed by Jaskot & Oey (2013) and Nakajima & Ouchi (2014) was successfully used to select LyC-leaking candidates but no strong correlation was found with the measured values of escape fraction (see Izotov et al. 2018b; Naidu et al. 2018; Bassett et al. 2019; Nakajima et al. 2020, and discussions therein). Based on a similar idea, the lack of emission from ions with low ionization potentials like [S II] λλ6716,6731 Å has also been proposed to target leaking candidates (Wang et al. 2019, 2021; Katz et al. 2020). This lack of emission of some low ionization species was first interpreted as the signature of a density-bounded galaxy where the outer part of H II regions were completely stripped out. However, using simple photoionization models, Stasińska et al. (2015) have shown that on average galaxies with high O32 cannot have massive escapes of ionizing photons, since low ionization lines like [O I]6300 Å are often also detected in these galaxies, implying the presence of radiation-bounded regions. Subsequently Plat et al. (2019) and Ramambason et al. (2020) noted that several strong LyC-emitters show surprisingly strong [O I]6300 Å emission, and proposed several explanations. While Stasińska et al. (2015) and Plat et al. (2019) suggested that such emission could be powered by the presence of AGN or radiative-shocks, we proposed in Ramambason et al. (2020) a 2-component model combining both density- and ionization-bounded regions.
Citation: Nakajima & Ouchi (2014)
Citation function: Uses; Uses
We emphasize that 7Be/H remains larger by at least one order of magnitude than predicted by nova models (Starrfield et al. 1978; Hernanz et al. 1996; José & Hernanz 1998). We note, however, that the final amount of 7Be is sensitive to the amount of 3He in the donor star as a higher abundance of 3He is expected to result in a higher 7Be abundance. Boffin et al. (1993) and Hernanz et al. (1996) found a logarithmic dependence of the 7Be output to the initial 3He abundance. The non-linearity of 7Be yields results from 3He(3He, 2p)4He and its importance increases as the square of the initial 3He abundance. Therefore, it produces a leaking of the available 3He for the 3He(α, γ) 7Be, whose rate increases only linearly for the initial 3He abundance. For 3He enhancements up to 100 solar, Boffin et al. (1993) derive X(7Be)/X(7Be0) = 1 + 1.5log X(3He)/X(3He⊙), where 7Be0 is the 7Be final mass fraction obtained with a solar initial 3He mass fraction. From the theoretical point of view, it is believed that low-mass main-sequence stars synthesize 3He through the p–p chains with peak abundances of few 10−3 by number (Iben 1967). As the star ascends, the red giant branch convection dredges up 3He-enriched material to the surface, which is later expelled into the interstellar medium by wind or during the planetary phase. 3He is a particularly difficult element to measure. It can be measured in H ii regions by using measurements at a frequency of 8.665 GHz (i.e. 3.46 cm), which is emitted naturally by ionized 3He (Bania, Rood & Balser 2010; Balser & Bania 2018) or in the stellar atmospheres of hot stars (Geier et al. 2012). Surprisingly, interstellar medium observations indicate that there is far less of this element in the Galaxy than the current models predict. In order not to overproduce 3He in the course of chemical evolution, it has become customary to assume that some unknown 3He-destruction mechanism is at work in low-mass giants (Dearborn, Steigman & Tosi 1996; Galli et al. 1997; Chiappini, Renda & Matteucci 2002; Romano & Matteucci 2003). For instance, Charbonnel & Zahn (2007) suggested a thermohaline mixing during the red giant branch phase of low-mass stars. However, in a few planetary nebulae, 3He/H is found to be high at the level of 10−3, consistent with predictions from standard stellar models (Rood, Bania & Wilson 1992; Balser & Bania 2018).
Citation: Iben 1967
Citation function: Background
In this study, we examine the stellar populations of NGC 5053 using the UVIT, on-board the AstroSat. NGC 5053 is a galactic GC that lies in the northern constellation of Coma Berenices, having Galactic coordinates l = 335${_{.}^{\circ}}$70, b = 78${_{.}^{\circ}}$95. This cluster is located at a distance of 17.54 ± 0.23 kpc (Baumgardt & Vasiliev 2021) and the metallicity is estimated to be [Fe/H] = −2.27 (Harris 2010). It is, therefore, placed amongst the metal-poor galactic GCs. The tidal radius of this cluster has been estimated as rt = 15.2 ± 3.3 arcmin by de Boer et al. (2019). When it was discovered in 1784, it was not classified as a GC because of its appearance and structure; there was no densely packed nucleus, and the central region was resolvable (Herschel 1786). These observations placed this GC in a grey-area between globular and open clusters. It was only in 1928 that Baade first classified it as a GC owing to its high latitude, richness in faint stars and presence of variable stars (Baade 1928). The variable stars in this cluster include RR Lyrae variables and SX Phoenicis (SX Phe) stars (Sawyer 1946; Nemec, Mateo & Schombert 1995a; Nemec et al. 1995b; Nemec 2004). The cluster is known to have a HB that is extended predominantly towards the blue flank of the RR Lyrae instability strip, which is an archetypal feature in many metal-poor GCs (Sarajedini & Milone 1995). A rich population of BSSs has been identified in this cluster (Sarajedini & Milone 1995). The presence of relatively large number of BSSs in such low density GCs has been examined to propose alternate formation mechanisms of BSSs (Leonard & Fahlman 1991). The cluster has been examined in the UV regime by Schiavon et al. (2012) using photometric data from GALEX, and they have published a catalogue of the UV bright stars present in this cluster. The location of this cluster in the sky and its velocity have given rise to a debate on whether it belongs to the Milky Way or Sgr dSph (e.g. Law & Majewski 2010; Boberg, Friel & Vesperini 2015; Sbordone et al. 2015; Tang et al. 2018).
Citation: Baade 1928
Citation function: Background
Situated in the closest vicinity of their host stars (0.05 AU) and having no counterparts in our Solar system, ultra-hot Jupiters (Arcangeli et al. 2018; Bell & Cowan 2018; Parmentier et al. 2018) are ideal testbeds for studying the impact of 3D effects on high-resolution spectra. There are two important reasons for this. First, ultra-hot Jupiters are accessible objects to observe. Their short orbital periods (1–2 d) and hot, extended atmospheres make them perfect targets for transmission spectroscopy (Hoeijmakers et al. 2019; Von Essen et al. 2019; Ehrenreich et al. 2020; Borsa et al. 2021), emission spectroscopy (Evans et al. 2017; Arcangeli et al. 2018; Mikal-Evans et al. 2020) and phase-curve studies (Zhang et al. 2018; Bourrier et al. 2020b; Mansfield et al. 2020). Secondly, ultra-hot Jupiters display extreme variations across their atmospheres, because they are expected to become tidally locked soon after their formation (Rasio et al. 1996; Showman & Guillot 2002). As a result, their atmospheres virtually consist two different worlds: a permanently irradiated dayside and a permanently dark nightside. The scorching, cloud-free dayside (T ≳ 2500 K) nearly resembles a stellar photosphere, where most molecules are dissociated1 and metals become ionized (Parmentier et al. 2018; Hoeijmakers et al. 2019). On the other hand, the nightside is substantially cooler (T ≲ 1000 K) and may even serve as a stage for cloud formation (Helling et al. 2019; Ehrenreich et al. 2020). Ultra-hot Jupiters also exhibit large differences in their thermal structures: the dayside is expected to show strong thermal inversions (Haynes et al. 2015; Evans et al. 2017; Kreidberg et al. 2018; Pino et al. 2020; Yan et al. 2020), whereas nightside temperatures are expected to monotonically decrease with altitude. Furthermore, ultra-hot Jupiters feature strong winds in the order of 1–10 km s−1 (Tan & Komacek 2019), which arise as a result of the continuous day-night forcing. Many observational studies have measured Doppler shifts due to winds on ultra-hot Jupiters (Casasayas-Barris et al. 2019; Bourrier et al. 2020a; Cabot et al. 2020; Ehrenreich et al. 2020; Gibson et al. 2020; Hoeijmakers et al. 2020; Nugroho et al. 2020; Stangret et al. 2020; Borsa et al. 2021; Kesseli & Snellen 2021; Rainer et al. 2021; Tabernero et al. 2021), yet inferring the underlying 3D circulation pattern is a formidable challenge.
Citation: Borsa et al. 2021
Citation function: Motivation
Therefore, the kinetics characteristics of the star could be the only piece to judge whether or not the SM star is the surviving companion of SN 1006. If its space velocity is significantly different from the other stars in the remnant of SN 1006, the probability to be the surviving companion would become high. Otherwise, the probability becomes low. We check the proper motion of the stars within 5 arcmin of the remnant centre from Gaia DR2, as shown in Fig. 19. From the figure, it seems that there is not difference between the SM star and other stars in the remnant of SN 1006 in the aspect of proper motion, i.e. the proper motion of the SM star only slightly deviates from the median value of the proper motions of the stars at the direction of the SNR centre of SN 1006, and such a proper motion disfavours the SM star as the surviving companion of SN 1006 (Schweizer & Middleditch 1980; Burleigh et al. 2000). So, a 3D space velocity is helpful to judge the nature of the SM star. However, unfortunately, some data of the SM star in Gaia DR2 are so uncertain that we cannot use them to constrain its 3D space velocity, otherwise we could obtain a complete wrong conclusion.5 For example, the parallax of the SM star is ϖ = 0.0736 ± 0.1244, and then σϖ/ϖ = 1.69 which is much larger than the threshold value of 0.2 for distance estimation from GAIA DR2 data (Astraatmadja & Bailer-Jones 2016; Katz et al. 2018). The distance of the SM star from this parallax is much larger than all the previous measurements from spectrum by at least a factor of 2 (see summary in Burleigh et al. 2000). Considering that some other astrometric measurements of the SM star are also very uncertain, we applied the measurements in the previous literatures as the distance of the SM star. Based on a radial velocity of $-13\pm 17\, {\rm km^{\rm -1}}$ and a distance of 2.07 ± 0.18 kpc (Schweizer & Middleditch 1980; Winkler et al. 2003; Kerzendorf et al. 2018), we can calculate the UVW velocities of the SM star, i. e. $U=-5.2\pm 14\, {\rm km^{\rm -1}}$, $V=197\pm 10\, {\rm km^{\rm -1}}$, and $W=3.1\pm 5\, {\rm km^{\rm -1}}$. The V value of the SM star is smaller than that of a normal disc star. We then transform these velocities into the Galactic rotational velocity at a Galactocentric distance of ∼6.67 kpc, i.e. $V_{\rm c}=196\pm 12\, {\rm km^{\rm -1}}$, which is smaller than the Galactic rotational velocity of the disc stars at the Galactocentric distance by $50\pm 19\, {\rm km^{\rm -1}}$ (Huang et al. 2016). This velocity difference is marginally consistent with the predicted orbital velocity here (see Fig. 7). In addition, the smaller rotational velocity of the SM star may explain its small proper motion shown in Fig. 19. So, the SM star is still possible to be the surviving companion of SN 1006.
Citation: Schweizer & Middleditch 1980
Citation function: Similarities
Citation: Schweizer & Middleditch 1980
Citation function: Uses; Uses
Active galactic nuclei (AGN) are thought to be powered by an accretion disc around a supermassive black hole, mostly emitting in the optical/ultraviolet (UV) band. According to the standard paradigm, the X-ray emission is due to thermal Comptonization of the soft disc photons in a hot region, the so-called corona (Haardt & Maraschi 1991; Haardt, Maraschi & Ghisellini 1994, 1997). This process explains the power-law shape of the observed X-ray spectrum of AGN. A feature of thermal Comptonization is a high-energy cut-off, which has been observed around ∼100 keV in several sources thanks to past observations with Compton Gamma-Ray Observatory (CGRO)/Oriented Scintillation Spectrometer Experiment (OSSE; Zdziarski, Poutanen & Johnson 2000), BeppoSAX (Perola et al. 2002) and more recently Swift/Burst Alert Telescope (BAT; Baumgartner et al. 2013) and INTEGRAL (Malizia et al. 2014; Lubiński et al. 2016). From the application of Comptonization models, such high-energy data allow to constrain the plasma temperature, which is commonly found to range from 50 to 100 keV (e.g. Zdziarski et al. 2000; Lubiński et al. 2016). Furthermore, the cut-off energy is now well constrained in an increasing number of sources thanks to the unprecedented sensitivity of NuSTAR up to ∼80 keV (e.g. Ballantyne et al. 2014; Brenneman et al. 2014; Marinucci et al. 2014; Baloković et al. 2015; Matt et al. 2015; Ursini et al. 2015). The primary X-ray emission can be modified by different processes, such as absorption from neutral or ionized gas (the so-called warm absorber), and Compton reflection from the disc (e.g. George & Fabian 1991; Matt, Perola & Piro 1991) or from more distant material, like the molecular torus at pc scales (e.g. Matt, Guainazzi & Maiolino 2003). A smooth rise below 1–2 keV above the extrapolated high-energy power law is commonly observed in the spectra of AGN (see e.g. Bianchi et al. 2009). The origin of this so-called soft excess is uncertain (see e.g. Done et al. 2012). Ionized reflection is able to explain the soft excess in some sources (e.g. Crummy et al. 2006; Ponti et al. 2006; Walton et al. 2013), while Comptonization in a ‘warm’ region is favoured in other cases (e.g. Mehdipour et al. 2011; Boissay et al. 2014).
Citation: Zdziarski et al. 2000
Citation function: Background