adsabs/WIESP2022-NER huggingface dataset:
Paper: 2017ApJ...844...60E (fulltext)
that ongoing accretion is occurring at such a slow rate that it is not detectable and that we caught H α in emission when the accretion was barely high enough to detect. Despite the evidence for some gas in the disk and its young age, the disk fractional luminosity and low dust mass of T 43 is consistent with what is seen in debris disks. We conclude that T 54 has ended its primordial disk lifetime and is crossing over the bridge to the debris disks stage. The fact that this is a binary system hastens its evolution, allowing us a rare glimpse of this bridge in a young star-forming region. Future observations of this object will help us understand this short-lived stage. Appendix T 54 has a nearby FIR source that contaminates the observed FIR emission (left, Figure 9 ; Matrà et al. 2012 ). Here we show that this unidentified FIR source should have no significant contribution to most of the data that we study in this paper. The unidentified FIR source is 6′′ away from T 54. Assuming it is in Cha I (at 160 pc), it is located 1040 au away from T 54 at a position angle of 196°. To the best of our knowledge, this FIR source does not overlap with previously reported objects. Figure 9. Left: PACS 70 μ m image from the HSA archive with T 54 and a nearby 2MASS source labeled with red triangles (top and lower, respectively). The FIR source identified by Matrà et al. ( 2012 ) is denoted with a black ellipse. Center: Spitzer IRS spectral footprint overlaid on the PACS 70 μ m image. We show the short-low (yellow) and long-low (black) slits as well as the two nods of the short-high observations (red). Right: Herschel PACS spectral footprint (red boxes) corresponding to the [O i ] observations discussed in Section 4.1.4 overlaid on the PACS 70 μ m image. Note that T 54 is in the central spaxel, which was used to extract the spectrum, and the unidentified FIR source is not covered by the central spaxel. Matrà et al. ( 2012 ) find that the unidentified FIR source is not detected in JHK -band photometry, IRAC photometry, or MIPS data taken at 24 μ m. However, it is detected in PACS 70 and 100 μ m data as well as in SPIRE photometry taken at 160, 250, and 350 μ m. The flux of this unidentified FIR source is comparable to or greater than the flux of T 54 at 100 μ m and beyond (Matrà et al. 2012 ). Therefore we show in Figure 8 the fluxes reported by Matrà et al. ( 2012 ) as upper limits since these data include emission from both the T 54 system and the nearby unidentified FIR source. It is unlikely that the unidentified FIR source is contributing to the Spitzer IRS spectra of T 54. As reported by Matrà et al. ( 2012 ), the unidentified FIR source is not detected in NIR, MIR IRAC, or MIPS images. To explore this further, we overlaid the IRS slit positions on a PACS 70 μ m image of T 54 (center panel, Figure 9 ). Low- and high-resolution IRS spectral observations of T 54 are available in CASSIS. In the low-resolution observations (2005 April 24 10:57:08), the
Paper: 2015MNRAS.454.2657E (acknowledgments)
We thank the referee for an encouraging and incisive report which improved the quality of this paper. We gratefully thank Thierry Foglizzo for fruitful exchanges and the mpi-amrvac development team, in particular Oliver Porth, Rony Keppens and Zakaria Meliani, for useful assistance. We acknowledge the financial support from the UnivEarthS Labex programme of Sorbonne Paris Cité (ANR-10-LABX-0023 and ANR-11-IDEX-0005-02). IEM acknowledges financial support from the ENS of Cachan through a contrat doctoral spécifique pour normaliens and from the University Paris 7 Diderot. The numerical simulations were run on the François Arago Centre facilities in Paris 7 and using the HPC resources from GENCI-CINES (Grant c2015047469).
Paper: 2015ApJ...806..137R (acknowledgments)
We would like to thank Kevin France for discussions concerning the MUSCLES database and Sukrit Ranjan for discussions concerning DNA action specta. This work has made use of the MUSCLES M dwarf UV radiation field database. We would also like to acknowledge support from DFG funding ENP KA 3142/1-1 and the Simons Foundation (290357, Kaltenegger and 290360, Sasselov).
Paper: 2016MNRAS.457..198S (fulltext)
can be asymmetric both in position (i.e. different distances from disc centre) and brightness (e.g. Baptista, Harlaftis Steeghs 2000b ; Baptista et al. 2005 ). Aside of the transient spiral structures seen during dwarf nova outbursts, spiral shocks were also observed in nova-like systems – where they presumably are persistent, tidally induced effects in permanently large and hot accretion discs (Baptista Bortoletto 2008 ; Saito et al. 2010 ). V348 Puppis is the optical counterpart of the faint HEAO X-ray source 1H0709-36, identified by Tuohy et al. ( 1990 ) as an eclipsing CV inside the period gap ( P orb = 2.44 h, Baptista et al. 1996b ; Rolfe, Haswell Patterson 2000 ). Its light curve, which shows the gradual eclipse of an extended source, and the observed rotational disturbance in its emission lines during eclipse indicate the presence of a bright, high mass accretion rate disc outshining the contributions from the WD at disc centre and the bright spot at disc rim. The detection as a hard X-ray source, the presence of a strong He ii λ4686 emission line in the spectrum and the possible existence of a second periodicity close to the orbital period led to the suggestion that V348 Pup might be a magnetic CV of the DQ Her type close to synchronize (Tuohy et al. 1990 ; Rosen et al. 1994 ). However, its spectrum exhibits strong single-peaked Balmer and He i emission lines that remain largely unobscured during disc eclipse (indicating significant contribution from above the orbital plane) and displays absorption components at phases opposite to the disc eclipse (φ = 0.4–0.5). This spectroscopic behaviour led Rodríguez-Gil et al. ( 2001 ) to classify V348 Pup as a nova-like variable of the SW Sex type. Moreover, the nature of the second periodicity became clear with the discovery of long-lasting, positive superhumps (with a period 6 per cent longer than P orb ), resulting from the gravitational interaction between a slowly precessing non-axisymmetric accretion disc and the mass-donor star (Rolfe et al. 2000 ). Its eclipsing nature and the location within the period gap make V348 Pup an obvious target for an investigation of its accretion disc properties with eclipse mapping techniques (Horne 1985 ). For this purpose, time series of UV and optical spectroscopy of V348 Pup were collected with the Hubble Space Telescope ( HST ). A first analysis of this data set were reported by Froning, Long Baptista ( 2003 ). They found that the HST out-of-eclipse spectrum is characterized by a blue continuum with strong emission lines and a broad dip in the spectral range λ2000–3000 Å (attributed to a blend of Fe ii absorption lines known as the ‘Fe ii curtain’; Horne et al. 1994 ), but the Balmer jump is shallow or absent. The eclipses are remarkably shallow across the observed spectral range (≃40–50 per cent of the out-of-eclipse light remains visible at mid-eclipse in the continuum, and up to 80 per cent in the emission lines) indicating that a significant
Paper: 2019MNRAS.485L..58E (fulltext)
that the progenitor was enshrouded in a dusty cocoon of material that was lost close to the explosion. If the SN was not an electron capture SN it would have been the explosion of a 11 M ⊙ red supergiant. However 2008S has little in common with other such SN from progenitors in this mass range (Smartt 2015 ; Davies Beasor 2018 ). In conclusion using the new distance to determine the luminosity of the progenitor we find it more likely than before that SN 2008S was an electron capture event. It also was likely to have been the result of a binary interaction which was suggested by Podsiadlowski et al. ( 2004 ) to be the most likely channel for such core-collapse events. While SN 2008S is at the minimum mass for a core-collapse SNe, N6946-BH1 is at the maximum mass for a core-collapse of a red supergiant. This star was observed to rapidly fade with a light curve that was in agreement with a source powered by a feeding black hole (Adams et al. 2017 ). When the nature of the progenitor was examined this study assumed a distance of 5.95 Mpc. Therefore again this increases the suggested progenitor luminosity from log ( L/L ⊙ ) = 5.3 to 5.5. As shown in Fig. 1 this increase in the luminosity implied a mass for the progenitor star or 24 M ⊙ that is above the predicted mass limit for the formation of a black hole. This is also consistent with the mass estimated from the surrounding stellar population by Murphy et al. ( 2018 ). Again while it is only one event tension between the mass estimates for black hole formation and the luminosity of the observed progenitor stars are reduced by this study. We note that this result is in agreement with the work of Murphy et al. ( 2018 ). 4 CONCLUSIONS We have re-evaluated the distance to NGC 6946. The value of 5.9 Mpc is incorrect and should not be used in future studies. We used the observed SN rate to estimate the star formation rate and thus the FUV luminosity of the galaxy. Using the standard candel method we derive a distance of 7.9 ± 4.0 Mpc. This is consistent with the more recent distances estimated by Anand et al. ( 2018 ) and Murphy et al. ( 2018 ) of 7.72 ± 0.32 Mpc and 7.83 ± 0.29 Mpc, respectively. We strongly recommend that all future studies concerning NGC 6946 use the updated distance of 7.72 ± 0.32 Mpc from Anand et al. ( 2018 ). Their distance is consistent with the distance of Murphy et al. ( 2018 ) but Anand et al. ( 2018 ) took additional steps to reduce contamination and is therefore preferred. We re-evaluated the nature of the five progenitors of observed core collapse events in NGC 6946 using the distance of 7.72 ± 0.32 Mpc. We found that the mass constraints change significantly. The two most significant changes are: first, that the observed progenitor of SN 2008S is now more consistent with predictions for the progenitors of electron capture SNe. We suggest that SN 2008S was an electron capture SN with the progenitor having experienced a binary interaction before explosion. Second, the
Paper: 2021MNRAS.502.1753T (fulltext)
the microlensing detection (e.g. Alcock et al. 2000 ), halo white dwarfs have been suggested as natural candidates to contribute to the dark matter content of the Galaxy (e.g. Oppenheimer et al. 2001 ). An intense debate arose on this issue, although later studies demonstrated that the white dwarf contribution to dark matter was rather limited (e.g. Torres et al. 2002 ; Flynn, Holopainen Holmberg 2003 ; García-Berro et al. 2004 ; Kilic et al. 2004 ; Bergeron et al. 2005 ). In any case, the search for white dwarfs in the Galactic halo has proven to be a difficult endeavour. In this sense, the intrinsic faintness and the low space density of halo white dwarfs, along with their high surface gravity, which erases any trace on metal content from their atmospheres and disables accurate radial velocity measurements due to the broadening of the Balmer lines, are some of the main factors that have hindered the discovery of suitable candidates. Consequently, during decades, the identification of halo members has relied on the search of cool objects in high proper motion surveys (e.g. Monet et al. 2000 ; Nelson et al. 2002 ; Hall et al. 2008 ; Catalán et al. 2012 ; Munn et al. 2016 ). It was not until the ESO SNIa Progenitor surveY project (see Napiwotzki et al. 2020 , and references therein) that radial velocities were measured for the first time with enough precision to recover the 3D kinematics of white dwarfs. That permitted to identify high eccentric retrograde orbits as strongly indicators of their belonging to the halo population (Pauli et al. 2006 ). Although the number of halo members was sparse, an estimate of the age of the inner halo has been possible ( Gyr; Kalirai 2012 ). A key ingredient of these studies is the white dwarf luminosity function. Defined as the number of white dwarfs per bolometric magnitude unit and cubic parsec, the white dwarf luminosity function is a valuable tool to derive the age, history, and evolution of the components of our Galaxy. Moreover, the luminosity function holds all the information about the cooling process of the white dwarfs, being then an excellent tool for testing the physics of evolutionary models (see García-Berro Oswalt 2016 , for a comprehensive review about these issues). However, the effectiveness of the white dwarf luminosity function requires the existence of a complete, well-defined volume sample. In this sense, the lack of accurate distances previous to the Gaia era, joint to the selection biases inherited from magnitude-limited samples, have hampered the achievement of more conclusive results in such an elusive sample as the halo white dwarf population. The first observational halo white dwarf luminosity function (Liebert et al. 1989 ) contained only five objects, and no more than 20 were proposed as halo candidates through a neural network artificial intelligent classification algorithm (Torres, García-Berro Isern 1998 ). Based on the SuperCOSMOS Sky Survey observations towards the southern
Paper: 2015MNRAS.454.1711B (acknowledgments)
We thank Paz Beniamini, Petar Mimica, Enrico Ramirez-Ruiz, Lorenzo Sironi, and Alexander van der Horst for useful comments on the manuscript. We acknowledge support from NASA grant no. NNX13AP13G.
Paper: 2021ApJ...908L..29G (fulltext)
15 Myr. This excess in BH number can be partially explained by an increase in binary interaction rates in these models. Specifically, model 2a has a higher number of binary mergers involving stars with zero-age main sequence (ZAMS) masses between 10 and 25 M ⊙ that would not have otherwise been BH progenitors (at 0.1 Z ⊙ , the minimum ZAMS mass that will form a BH is roughly 20 M ⊙ ). However, these merger products do evolve into BHs with masses roughly in the range of M BH ≈ 5–15 M ⊙ . For models 2e and 2k , the higher number of BHs can be partially attributed to mass transfer in a binary increasing the mass of one of the components past the threshold of BH formation. Naturally, some of these differences are clearly due to stochastic fluctuations in the initial evolution of the model, leading to divergent paths in the dynamical evolution of the system. The main takeaway is that binary evolution processes, coupled with dynamics, can greatly influence BH growth. We leave a detailed study of the extent to which binary evolution processes affect the long-term cluster evolution and dynamics for a future study. It is important to note that many (roughly 55%) of these “excess” BHs formed with masses less than 15 M ⊙ are ejected promptly from their host cluster through natal kicks. Low-mass BHs are expected to form with less mass fallback and thus receive larger natal kicks (e.g., Fryer et al. 2012 ). Thus, this excess of BHs formed is unlikely to have a significant effect on the long-term cluster dynamics. 4. Discussion and Conclusions Our results show that high binary fractions for massive stars lead to an increase in binary-mediated dynamical interactions, which in turn have an important effect on massive BH formation. We have demonstrated that increasing the high-mass binary fraction, consistent with observations, while keeping all other cluster parameters (e.g., cluster masses, virial radii, metallicity) fixed at the values used in Kremer et al. ( 2020b ), may significantly change the early evolution of YMCs. In particular, high-mass binaries facilitate high rates of massive star collisions (occurring through both dynamical encounters and binary-evolution-driven mergers) that can lead to the formation of massive BHs, both within and above the pair-instability mass gap. With the exception of the high-mass binary fraction, the parameters for the models calculated in this Letter are identical to those in Kremer et al. ( 2020b ), which were shown to lead to present-day properties matching well those of Milky Way GCs. If indeed the majority of high-mass stars have binary companions at birth, massive BH formation may therefore be a common occurrence in the early phases of GC evolution. As shown in Figure 2 , the formation of massive BHs is sensitive to small changes in initial virial radius. This is in agreement with previous works on the topic of IMBH formation that have shown that the onset of a collisional runaway occurs abruptly under small
Paper: 2018MNRAS.480.2357K (fulltext)
index Γ and the inner disc temperature are consistent with the results of previous studies of M33 X-8, where a similar spectral model and data from Swift /XRT (La Parola et al. 2015 ), XMM–Newton (Foschini et al. 2004 ), or Suzaku (Isobe et al. 2012 ) were used. The innermost disc radius, R in , can be directly estimated from the normalization of the diskbb component: , where D is the distance to the source in units of 10 kpc and i is the inclination of the disc. The estimated inner disc radius, R in = (37 ± 2)cos −1/2 i km, is in good agreement with the results of previous studies (e.g. Middleton et al. 2011 ). Associating this size with the radius of the innermost stable orbit of a standard accretion disc suggests that the relativistic compact object has a mass of at least ∼4 M ⊙ , with the exact value depending on the (unknown) inclination of the disc. More accurate estimation should take into account the inner boundary condition, the spectral hardening and relativistic corrections, which in turn would depend on the structure of the accretion disc (whether it is thin or slim). Nevertheless, the above crude estimate is consistent with the compact object being a stellar-mass black hole. Replacing the power-law component by its modification with an exponential cut-off (Model 5) further improves the fit statistics to χ 2 /d.o.f. =891/900, indicating the presence of a high-energy cut-off in the spectrum (the F -test probability is 10 −14 for the null-hypothesis Model 4). We conclude that the broad-band 0.3–20 keV spectrum of M33 X-8 cannot be described by a single accretion disc model diskpbb , which was previously used to successfully fit some narrow-band spectra of this source (e.g. La Parola et al. 2015 ). The broad-band spectrum strongly requires an additional power-law component with a possible high-energy cut-off. 3.3 Physically motivated models We have demonstrated that the broad-band spectrum of M33 X-8 certainly contains two components: the first one is well approximated by a multicolour disc with the inner temperature of kT ∼ 1 keV, and the second one is a power-law-like high-energy tail detected up to 20 keV. As often discussed in the literature (e.g. Gierliński et al. 1999 ; Steiner et al. 2009 ), the combination of a low-energy blackbody-like component with a high-energy power law does not allow one to properly model the spectrum at low energies, because a steep power law extending to the soft part of the spectrum requires the introduction of an artificially high absorption column density and can result in erroneous determination of the disc properties. This problem can be solved by using a self-consistent Comptonization model that has a low-energy cut-off in the otherwise power-law-like spectrum, such as the convolution model simpl (Steiner et al. 2009 ), which provides a description of Comptonization based on the non-relativistic theory developed by Sunyaev Titarchuk ( 1980 ). In this model, a fraction of the photons from an
Paper: 2016A&A...585A..51D (fulltext)
images are for the HST bands and for the Spitzer bands. We adopt a -CDM cosmological model with km s Mpc, , and . We assume a Salpeter IMF ( Salpeter 1955 ). All magnitudes are expressed in the AB system ( Oke Gunn 1983 ). 2. Data 2.1. Spectroscopy The available low-resolution (LR, ) and medium-resolution (MR, ) VLT/VIMOS spectroscopy covering the spectral ranges 3400–7000 Å and 4800–10 000 Å were obtained during 2005, with 4hr exposure time for both LR and MR spectra. The LR spectrum continuum sensitivity is at 3700 Å with 4 h and assuming seeing and airmass 1.2. It is calculated over one pixel along dispersion and (full width half maximum) along spatial direction. ( Popesso et al. 2009 ; Balestra et al. 2010 ). Slit widths of were adopted and the median seeing during observations was . We have newly reduced both configurations, in particular focusing on the Lyman continuum emission (LR-spectrum) and the ultraviolet spectral properties in the non-ionizing domain (MR-spectrum). Reduction was done in two ways: as described in Balestra et al. (2010) and the AB–BA method as described in Vanzella et al. (2011 , 2014b) . Both methods produce consistent results. Recently, a near infrared spectrum of Ion2 was obtained with Keck/MOSFIRE in March 2015 (PI: Günther Hasinger). Integration time of 32, 32, and 36 min were obtained in the -, -, and -bands, respectively, with the AB–BA dithering pattern of 2.9. The adopted slit width was 0.7 producing a dispersion of 1.30 Å, 1.63 Å, and 2.17 Å per pixel, in the -, -, and -bands, respectively; the pixels scale along the spatial direction was 0.18/pixel. Ion2 was inserted in a mask aimed to target low- objects as well. The -band is useful for monitoring low- emission lines. Reduction was performed using the standard pipeline 1 . Briefly, two-dimensional sky-subtracted spectra, rectified and wavelength calibrated, are produced from which the one-dimensional variance weighted spectrum is extracted with a top-hat aperture. The wavelength solution was computed from sky emission lines (e.g., Rousselot et al. 2000 ) in the three bands reaching an rms accuracy better than 1/20, 1/30, and 1/40 of the pixel size in -, -, and -bands, respectively (typical rms 0.05 Å). Table 1 Emission line properties. Notes. (a) Observed emission line flux and error corrected for slit losses (50%) in units of erg s cm. In parenthesis, the [O iii ] flux adopting the theoretical ratio ( Storey Zeippen 2000 ). (b) The emission line dust correction method is described in Appendix A . (c) FWHM (instrumental corrected) and error in units of km s. (d) Equivalent width and error in units of Å. (e) Sum of the two components. (f) Main component. (g) Marginal detections (Fig. 6 ): , , and . (h) FWHM is the contribution of the two blended components, 1907–1909 Å. (i) The error is large because the line is on the top of a sky line (Fig. 4 ). Flux calibration was performed by observing the standard star A0V HIP 26577. Given the slit width () and the