adsabs/WIESP2022-NER huggingface dataset:
Paper: 2019MNRAS.483..332G (fulltext)
studied the typical outcomes of close encounters between planets and orbit-crossing planetesimals. Applying their results to the fifth planet shows that it would scatter most of the orbit-crossing material on to bound orbits, thus supporting the fifth planet as the cause of the scattered disc. However, this planet’s mass of is much lower than the mass of the scattered population (). It is impossible that a much less massive planet scattered the much more massive planetesimal population. Read et al. ( 2018 ) furthermore found that planets with a mass in the range from to corresponding to either a semimajor axis from 115 to 130 au or from 140 to 160 au, with the most massive planets being nearest to or farthest away from the star, similarly well explained the observations. Although any planet with a mass in excess of would eject orbit-crossing material, a fifth planet capable of scattering of planetesimals is still possible. As the fifth planet may or may not be real, we shift our focus on to the four confirmed planets. Referring again to the work by Wyatt et al. ( 2017 ), one sees that the gas giants of HR 8799, each for itself, would typically eject their orbit-crossing material in unbound orbits. This material would be on its way out of the system, but, unless this is a continuous process, it should only operate for a relatively short amount of time. For this process to be continuous, a constant flux of material would need to cross the giant planet’s path, which is unlikely. Therefore, we do not consider scattering through the giant planets to be responsible for the scattered disc either. Migration is considered to have occurred in the formation history of HR 8799, as neither core accretion nor gravitational instability could have produced all four planets at their current location (Currie et al. 2011 ). Although we focus in the discussion on outward migration, a similar case can be made for inward migration. Using the formula of Ida et al. ( 2000 ) to calculate the migration rate in our primordial disc, we find a rate of a few 10 s of . Such a fast migration would allow the planets to quickly reach their current location. Taking into account that migration stops when the mass within a few Hill’s radii becomes less than the planet’s mass (Kirsh et al. 2009 ), we find the fully formed planets in HR 8799 to be too massive to have undergone migration within the protoplanetary disc. However, if the planets accreted their gas envelope while migrating or even after the migration ceased, then the core masses rather than planet masses would have to be considered. Migration rate of these, less massive, cores would be lower, and they could easily create a scattered disc of planetesimals instead of ejecting them out of the system. Planet–planet scattering may have also occurred early on in the history of the HR 8799 system. If the planet cores formed closer to the star by core accretion, such a scattering event could have relocated the planets, bringing
Paper: 2016MNRAS.463.3409V (acknowledgments)
We are grateful to M. Beasley, J. Beckman, F. La Barbera and R. Peletier for very useful discussions and valuable suggestions. We thank G. Bruzual for providing us with various versions of his models in the UV spectral range and for interesting suggestions. We also thank E. Toloba for providing us with fully reduced early-type galaxy spectra. We are grateful to the anonymous referee for valuable suggestions that helped us to improve the paper. This research has made an extensive use of the SIMBAD data base and VizieR catalogue access tool (both operated at CDS, Strasbourg, France), the NASA's Astrophysics Data System Article Service. This work has been supported by the grants AYA2013-48226-C3-1-P and AYA2013-48226-C3-2-P from the Spanish Ministry of Economy and Competitiveness ( MINECO ) and the Generalitat Valenciana under grant PROMETEOII/2014/069 . ER acknowledges a Marie Heim-Vögtlin grant from the Swiss National Science Foundation .
Paper: 2018MNRAS.477.1958C (acknowledgments)
ACKNOWLEDGEMENTS We wish to thank Michael Opitsch for useful discussion during the preliminary phase of this work. LC thanks the Department of Physics and Astronomy of the Padova University for hospitality while this paper was in progress. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.
Paper: 2020ApJ...900..152Z (fulltext)
starbursts that have inclination-corrected central brightness ( = 22.6 mag arcsec −2 for a kinematic inclination = 452) comparable to VCC 848 (Lelli et al. 2014 ). Because the rotation velocity gradient slope is proportional to the square root of local matter density for a solid-body rotation curve, we can infer that the merging event of VCC 848 has not built up an exceptionally compact core. 4.3. Dynamical and Baryonic Mass Profiles We use the rotation curve of the approaching half to determine the cumulative dynamical mass profiles (=/G), which are shown in Figure 5 . We also overplot mass profiles of the baryonic components, including stars and atomic and molecular gas, in Figure 5 . All of the profiles are extracted using the same geometric parameters as determined from the 2 dbat kinematic fitting. The stellar mass is estimated by using the ( g − i )–mass-to-light relation calibrated for Local Group dwarf galaxies (Zhang et al. 2017 ). Details about the procedure of stellar mass estimation are given in Paper I . We adopt a logarithmic stellar mass uncertainty of 0.2 dex, which is applicable to color-based mass estimation for low-dust-extinction cases (Zhang et al. 2017 ). The atomic gas mass is equal to the H i gas mass multiplied by a factor of 1.36 to account for helium. The molecular gas mass is indirectly inferred from the SFR estimated with H α luminosities (Catalan-Torrecilla et al. 2015 , Equation (10)) by adopting a constant molecular gas consumption timescale of (5.25 ± 2.5) × 10 −10 yr −1 found for nearby disk galaxies at (sub)kiloparsec scales (Bigiel et al. 2008 ; Leroy et al. 2008 ). 18 18 Another method of estimating the molecular gas mass is based on a metallicity-dependent gas-to-dust ratio (G/D). According to the G/D−metallicity relation from Rémy-Ruyer et al. ( 2014 ), VCC 848 is expected to have a of 3.00 ± 0.37 for (O/H) = 8.03 (Table 1 ). In addition, Grossi et al. ( 2015 ) determined a total dust mass of 10 5.01 M ⊙ (the value has been adjusted to reflect our adopted distance) within the central ∼30′′ of VCC 848. Therefore, the total gas mass (atomic plus molecular) within the central 30′′ is estimated to be , and hence the most likely molecular gas mass is ∼10 7.5 M ⊙ , found by subtracting the atomic gas contribution (Figure 5 ). This molecular gas mass is in reasonable agreement with that (10 7.6 M ⊙ ) inferred from the SFR-based method shown in Figure 5 . We choose to use the SFR-based method in this paper. This is because, besides their low spatial resolution, the Herschel images of VCC 848 generally have signal-to-noise ratios that are too low (except for the brightest star-forming sites) to allow for constructing a robust radial profile of dust mass densities. Figure 5. Accumulated radial profiles of the dynamical mass as well as the stellar, atomic, and molecular gas masses within the central 30′′. The same geometric parameters derived from the 2 dbat fitting of the H i velocity field (Figure 3 ) are used to
Paper: 2018MNRAS.480.5447Z (fulltext)
2.2 GTC spectroscopy Candidates have been followed up with optical spectroscopy using the Optical System for Imaging and low-Resolution Integrated Spectroscopy (OSIRIS; Cepa et al. 2000 ) instrument on the Gran Telescopio Canarias (GTC) since 2012. Table 2 shows a summary of our GTC spectroscopic observations. The spectra were mostly obtained using the R500R grism, with three making use of the R300R grism. These provide resolving power of around 500 and 300, respectively. The spectra were reduced using standard procedures with IRAF 1 . The HgAr, Ne, and Xe arcs were used on the wavelength calibration of our spectra. Spectral flux calibrations were achieved with standard stars (Oke 1974 , 1990 ; Hamuy et al. 1994 ; Bohlin, Colina Finley 1995 ). Spectral features caused by telluric absorptions in these spectra were not corrected. Five objects were observed twice to gain additional exposure time, and for the purposes of this paper we combined the spectra of these objects from both epochs. Table 2. Summary of the characteristics of the spectroscopic observations made with the OSIRIS on the GTC. R500R and R300R grisms cover a wavelength range of 480–1020nm. A 0.8 arcsec slit used for all observations. Note . References for flux calibration standard stars are 1. Bohlin et al. ( 1995 ); 2. Hamuy et al. ( 1994 ); 3. Oke ( 1974 ); 4. Oke ( 1990 ). 2.3 VLT spectroscopy We have also obtained optical-NIR spectroscopy of five of our candidates, a known sdL8 subdwarf and three radial velocity standard L dwarfs using the X-shooter spectrograph (Vernet et al. 2011 ) on the Very Large Telescope (VLT) since 2014. Table 3 shows a summary of our VLT spectroscopic observations. These X-shooter spectra were observed in an ABBA nodding mode with a 1.2 arcsec slit providing a resolving power of 6700 in the VIS arm and 4000 in the NIR arm. These spectra were reduced with ESO Reflex in a lazy mode (Freudling et al. 2013 ). The flat field used for the science data was also used for flat-fielding the flux standard star spectrum (by set flat strategy to false). Telluric corrections were achieved for both VIS and NIR arms using telluric standards that were observed right after or before our targets at a very close airmass. To extract telluric spectra for correction, we fitted and normalized spectra of telluric standards stars to their continuum and removed non-telluric features (e.g. hydrogen absorption lines) with iraf splot . These X-shooter spectra were smoothed by 101 and 51 pixels in the VIS and NIR arms, respectively, for display in figures, which reduced the resolving power to about 600 and increased their signal-to-noise ratio (SNR per pixel) by 10 and 7 times. Table 3. Summary of the characteristics of the spectroscopic observations made with X-shooter. Wavelength ranges for the VIS and NIR arms are 530–1020 and 990–2480 nm. 1.2 arcsec slits are used for both VIS and NIR arms. DENIS-P J025503.3−470049, 2MASS J08354256−0819237, and DENIS-P
Paper: 2017MNRAS.469..416L (fulltext)
fraction of the kinetic energy transferred to CRs by a single SN and SNR is the comoving SN rate. In principle, one should account also for helium nuclei. For simplicity, we assume here that particles can be treated as four protons, and hence be absorbed in the proton spectrum efficiency. In addition to the SFR, to derive SNR we need to know the number of SNe explosions per solar mass of forming stars, which is given by where ( m ) is the initial mass function of Population II/I stars, for which we assume the following form: with x 1.35 and m c 0.35 M ; m lies in the range [0.1, 100] M (Larson 1998 ). We neglect the contribution from Population III stars. Their SF history is very debated and still highly uncertain. In fact, detailed studies, exploiting cosmological hydrodynamical simulations implementing chemical feedback effects, have shown that Population II/I stars dominate the global SFR at any redshift (Tornatore etal. 2007 ; Maio etal. 2010 ). Moreover, the recent surveys hunting for Population III stars in the Milky Way have found no metal-free stars so far, implying that they are rare in the Milky Way even if they exist (Beers Christlieb 2005 ). Combining the above formulae, equation ( 19 ) becomes Particle acceleration in SN explosions is believed to occur through diffusive shock acceleration, which leads to momentum power-law spectra of accelerated particles. With this in mind, we assume that the source function of volume-averaged CR protons injected by SNe (defined as a rate per unit energy and volume) is where E is the proton kinetic energy, E 0 1GeV, m p is the proton mass, v / c is the dimensionless velocity of the particle, 2 is the slope of the differential spectrum of accelerated particles and C ( z ) is a redshift-dependent normalization obtained by imposing that the total kinetic energy rate equals , i.e. In equation ( 24 ), we fix E max 10 6 GeV (Lagage Cesarsky 1983 ), and we verify a posteriori that our conclusions are not strongly dependent on our choice of E min 10keV. From equation ( 23 ), one can easily realize that is the key parameter determining the fraction of the total kinetic energy released that goes into protons with E 1GeV. We keep as the only free parameter of the model. In Fig. 3 , we plot the source function as a function of the kinetic energy for three different values of . We found that protons with kinetic energies below 10MeV represent 0.5percent, 2.7percent and 13percent of the total kinetic energy released for 2,2.2 and 2.5, respectively. Figure 3. Source function of CR protons with respect to their kinetic energy at z 10 for a spectrum slope 2 (blue line), 2.2 (green) and 2.5 (red). 3.2 Energy losses in the IGM CRs can be an efficient heat source especially for a low-density gas. When the CR proton ionizes an atom, it transfers a certain fraction of its kinetic energy to the electron, which is either used for further atomic excitation and ionization or distributed via elastic collisions to other
Paper: 2018A&A...610A..74K (acknowledgments)
Acknowledgments The Herschel spacecraft was designed, built, tested, and launched under a contract to ESA managed by the Herschel/Planck Project Team by an industrial consortium under the overall responsibility of the prime contractor Thales Alenia Space (Cannes), and including Astrium (Friedrichshafen) responsible for the payload module and for system testing at spacecraft level, Thales Alenia Space (Turin) responsible for the service module, and Astrium (Toulouse) responsible for the telescope, with in excess of a hundred subcontractors. SPIRE has been developed by a consortium of institutes led by Cardiff University (UK) and including Univ. Lethbridge (Canada); NAOC (China); CEA, LAM (France); IFSI, Univ. Padua (Italy); IAC (Spain); Stockholm Observatory (Sweden); Imperial College London, RAL, UCL-MSSL, UKATC, Univ. Sussex (UK); and Caltech, JPL, NHSC, Univ. Colorado (USA). This development has been supported by national funding agencies: CSA (Canada); NAOC (China); CEA, CNES, CNRS (France); ASI (Italy); MCINN (Spain); SNSB (Sweden); STFC, UKSA (UK); and NASA (USA). This paper has made use of HSpot HIPE, which are joint developments by the Herschel Science Ground Segment Consortium, consisting of ESA, the NASA Herschel Science Center, and the HIFI, PACS and SPIRE consortia. This paper has made use of up-to-date SMARTS optical/near-infrared light curves that are available at the url: www.astro.yale.edu/smarts/glast/home.php . This work was partly supported by the NCN grant No. 2013/09/B/ST9/00599.
Paper: 2017ApJ...836....6A (fulltext)
in a scattering process the frequencies of the incoming and outgoing photons are completely uncorrelated. Therefore, it is very important to clarify when this limit can be safely applied. As mentioned before, the hypothesis of CRD is generally suitable for treating weak spectral lines, and it is a good approximation for modeling the core region of strong lines of the intensity spectrum. On the other hand, it turns out to be completely unsuitable for modeling the extended wings of strong resonance lines. In light of this, we provide a detailed comparison between the results obtained through the full PRD calculation and in the limit of CRD. In Figure 3 , the emergent intensity and Q / I profiles calculated in the absence of a magnetic field are shown, for the Sr ii 4078 Å line and for the Sr i 4607 Å line, comparing the full PRD calculation (solid line) to the calculation in the CRD limit (dashed–dotted line). This limit has been obtained by setting the elastic collision rate , while keeping all other parameters, including the depolarizing collision rates, as they were. In other words, we artificially modify the branching ratios so that all scattering processes occur through the redistribution matrix. For both lines considered, the intensity is greater at line center when using the CRD approximation than when PRD is considered. This can be qualitatively understood as follows. In the limit of CRD, the scattered radiation at a given frequency depends on a weighted average of the incoming radiation over the whole line profile, and it thus takes into account that, for an absorption line, the intensity increases going from the core to the wings. On the other hand, in the limit of CS the scattered radiation is related only to the incoming radiation at that same frequency. PRD phenomena relax such coherency, relating the scattering radiation to the incoming one, but over an interval that is generally smaller than the one considered in the CRD case. This same reasoning explains why, in PRD, the intensity is instead larger in the wings. Figure 3. Intensity (top row) and Stokes Q / I (bottom row) profiles of the emergent radiation calculated for an LOS with in the absence of a magnetic field, for the Sr ii line at 4078 Å (left column) and for the Sr i line at 4607 Å (right column), considering the FAL-C atmospheric model. In each figure, the results for the full PRD calculation (solid line) and the results obtained in the CRD limit (dashed–dotted line) are plotted. The reference direction of positive Q / I is taken perpendicular to the vertical direction. Figure 3 clearly shows that the difference between CRD and PRD calculations is, as expected, rather small in the Sr i 4607 Å line. In particular, we expect that, once the Stokes profiles are smeared to properly account for the effects of the atmosphere’s macroturbulent velocity and of the finite spectral resolution of a typical instrument, the differences between both calculations will be negligible. The
Paper: 2019A&A...624L...8P (fulltext)
3 shows the measured fractional linear polarization as a function of wavelength for the six stars measured in the B , V , R , and I bands. Fixing K = 1.15 ( Serkowski et al. 1975 ) and leaving λ max and p max as free parameters, we fit the six polarization profiles to Eq. ( 1 ) using weighted least-squares. Fig. 3. Wavelength dependence of polarization for six of the stars in the R -band sample (USNO-B1 identifiers shown in the legend). Debiased fractional linear polarization ( p̂ ) measured in the Johnsons-Cousins B , V , R , and I bands (horizontal axis shows effective wavelength). For each star, a line shows the best-fit Serkowski law (Eq. ( 1 ), with K set to 1.15). The fit λ max are in the range 580 ± 46 nm–643 ± 34 nm, while the fit p max are in the range 0.9%−1.7%. Although limited by the uncertainties of the measurements, we find no evidence for substantial differences between the best-fit value of λ max in this sample compared to canonical values ( Serkowski et al. 1975 ). If grain alignment were unusually efficient in this region with smaller grains able to be aligned, we might have expected a shift of λ max toward shorter wavelengths than the typical 550 nm ( Martin et al. 1999 ). This is not observed. 3.3. Magnetic field morphology If the magnetic field morphology is responsible for the high submillimeter polarization fractions, then any region with high p 353 should have a magnetic field that lies mostly in the plane of the sky with no substantial change in orientation either along the line of sight or within the beam. The orientation of the magnetic field as projected on the plane of the sky is probed by the optical polarization angle θ R as well as the submillimeter polarization angle rotated by 90°. Within an 80′ beam centered on this region we find χ 353 + 90° =112.8 ± 0.6°. The polarization angles inferred from stellar polarimery do not vary significantly across the region (Fig. 4 , left). The standard deviation of the distribution of θ R is consistent with it arising from measurement uncertainties, suggesting little beam depolarization. Fig. 4. Left : distribution of stellar polarization angles, θ R for stars farther than 250 pc (solid line). The dashed vertical line shows the χ 353 +90° from Planck at 80′ resolution. Right : probability density functions of the orientation of linear HI structures ( θ RHT ) for different velocity ranges (shown on the right). We only consider velocity ranges within the peak of the HI intensity spectrum. The filled circles in the bottom mark the mean value, ⟨ θ RHT ⟩ for each velocity range. All angles are measured in the Galactic reference frame, according to the IAU convention. Linear structures seen in HI have also been shown to trace the interstellar magnetic field ( Clark et al. 2014 ). Recently, Clark (2018) proposed that changes in the orientation of these structures in velocity space can be used to probe variations of the magnetic field orientation along the line of sight. By
Paper: 2018ApJ...856...77X (acknowledgments)
The authors thank the anonymous referee for constructive comments that helped improve this paper. This work is supported by National Science Foundation of China through grants NSFC-11303015 and NSFC-11133004, the National Key Research and Development Program of China (2016YFA0400803), NSFC-11773015 and NSFC-11133001. Facility: CXO (ACIS) .
Paper: 2016A&A...587A..22K (acknowledgments)
This research was supported by GA ČR 13-10589S. This research has made use of the Washington Double Star Catalog maintained at the US Naval Observatory. The authors would like to thank the anonymous referee for his/her comments on the manuscript, which led to significant improvement of the paper.
Paper: 2019A&A...622A.136H (acknowledgments)
We thanks R. Angulo, C. Hernández-Monteagudo, and O. Hahn for useful discussions. G.H. and R.A, acknowledge support from Spanish Ministerio de Economía and Competitividad (MINECO) through grant number AYA2015-66211-C2-2. U.K. acknowledges support by the Israel Science Foundation (ISF grant No. 1769/15). We acknowledge the use of HEALPix ( Górski et al. 2005 ).