adsabs/WIESP2022-NER huggingface dataset:
Paper: 2017MNRAS.465.1645H (fulltext)
strength of the magnetic field (Gehman, Adams Watkins 1996b ). The purpose of this paper is to add the effect of the ambipolar diffusion in the gravitational instability of filamentary MCs with magnetic field using non-ideal MHD. The outline of this article is as follows: in Section 2 , we present the basic MHD equations, which include the induction equation with an ambipolar diffusion term. We also describe the physical assumptions that are applied to the system. The unperturbed state of the system is introduced in Section 3 . In Section 4 , we perform linear perturbation analysis and then describe briefly the boundary conditions (BCs) and numerical method which are used to solve the perturbed equations in Sections 5 and 6 , respectively. We express the results in Section 7 . In Section 8 , we discuss our findings and finally, we summarize our conclusions in Section 9 . 2 MHD EQUATIONS WITH AD If the degree of ionization in a medium is very low, the force associated with the pressure gradient of the ions and their gravitational force could be neglected in comparison with the Lorentz and ion-neutral drag force. Accordingly, the Lorentz force and the drag force exerted by the neutrals on the ions could virtually balance each other. The equations of MHD in this so-called strong coupling approximation (Shu 1983 ) with self-gravity and AD term are as follows: where is the neutral gas density, is the fluid velocity, p is the pressure, is the gravitational potential and is the magnetic field strength, whereas A is the AD coefficient indicated by (Shu, Adams Lizano 1970 ), where i and are the ions density and the drag coefficient, respectively. Draine, Roberge Dalgarno ( 1983 ) proposed the value of 3.510 13 cm 3 g 1 s 1 . Due to the cosmic radiation, MC are partially ionized. Elmegreen ( 1979 ) showed that one could approximate the relation between ions and neutrals density as where C is a constant equals to C 310 16 cm 3/2 g 1/2 . So the equation ( 5 ) can be recast as where is a substitution for C and equals to 10.510 3 cm 3/2 g 1/2 s 1 . These assumptions lead to a fractional ionization of n i / n n 10 7 for n n 10 4 cm 3 , that is completely centred within 10 6 to 10 8 of observation of MC cores. By making the equations dimensionless, and A become 11.465 and 0.007 3/2 , respectively (see Section A ). The aforementioned fluid equations, are closed by addition of one equation of state, which is considered here to be isothermal, i.e. and, where c s is the isothermal sound speed which is about 0.2kms 1 for an MC with temperature of 10K and mean molecular weight of 2.36. The above values for the number density and sound speed, will give the value of 0.7 for a typical MC threaded by a magnetic field of 10G. 3 UNPERTURBED STATE The unperturbed fluid is supposed to be in the hydrostatic equilibrium state ( = 0). We make the assumption that the initial magnetic field does not have any role in stability of the filament. Therefore the density and
Paper: 2018MNRAS.475.3086L (acknowledgments)
Acknowledgements C-HL is indebted to the referee, whose comments greatly improved the manuscript. Based on observations (GN-2017B-FT-4) obtained at the Gemini Observatory and processed using the Gemini iraf package, which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the National Research Council (Canada), CONICYT (Chile), Ministerio de Ciencia, Tecnología e Innovación Productiva (Argentina), and Ministério da Ciência, Tecnologia e Inovação (Brazil). The author wishes to recognize and acknowledge the very significant cultural role and reverence that the summit of Maunakea has always had within the indigenous Hawaiian community and is very fortunate to have the opportunity to conduct observations from this mountain.
Paper: 2019MNRAS.487.3288S (fulltext)
was larger than the current discovered value, increasing the critical migration rate. The rate of migration can slow as the planets approach the inner edge of the disc due to the decrease in surface density (Kley Haghighipour 2014 , 2015 ). This slowing makes the likelihood of scattering more probable since it increases the time of interaction with the N:1 resonances. Additionally, if there exists a lack of circumbinary planets and especially multiplanet systems, this might be indicative of slow migration rates. 6.3 Application to the Pluto and Charon binary With a mass ratio of 0.1184, the dynamics around Pluto and its largest moon, Charon, behaves similar to the binary star systems described above. The binary dwarf planet has four discovered circumbinary moons: Styx, Nix, Kerberos, and Hydra. These moons are near to the 3:1, 4:1, 5:1, and 6:1 resonances with Charon and are therefore close to resonances with each other. None of these resonances are currently active according to the best estimates for the orbits and masses of the moons (Showalter Hamilton 2015 ). The very low eccentricities of all the bodies in the system makes active resonances unlikely due to the small resonance widths. Showalter Hamilton ( 2015 ) suggests the masses of the circumbinary moons are larger than predicted by albedo alone (Brozović et al. 2015 ) due to a librating three-body resonant angle for higher masses. While moon–moon resonances may help constrain the mass of the moons, the location of the moons near the N:1 resonances remains unexplained. The very low eccentricity of the moons suggests that in situ formation is likely (Brozović et al. 2015 ) but the collision from which the binary formed is not thought to have created enough leftover material at large enough distance to form the moons at their current location (Canup 2005 , 2011 ). Bromley Kenyon ( 2015 ) demonstrate that processes like physical collisions and collective gravitational scattering can expand the debris in the system to larger distances, facilitating formation. Resonant transport cannot explain the migration of all moons simultaneously since different N:1 resonances require different binary eccentricity for migration (Lithwick Wu 2008 ; Cheng et al. 2014 ). Resonant transport due to an eccentric Charon would also excite the eccentricity of the moons. Even outside of a resonance, the secular effects can excite the eccentricity of the moons above their current values. Formation models that predict a high initial eccentricity for the binary need to explain the current low eccentricity of the moons and overcome the likelihood of instabilities from the number of resonances in an eccentric system. Previous efforts have used stability criteria to constrain the maximum eccentricity of Charon as well as the maximum mass of the moons (Youdin, Kratter Kenyon 2012 ; Smullen Kratter 2017 ). The instabilities described here can also limit the possible masses and formation history of the moons. The potential
Paper: 2021MNRAS.502.5508P (fulltext)
more massive galaxies, LTGs have the same negative gradient regardless of mass. Metallicity gradients steepening for more massive ETGs is consistent with other studies from the MaNGA survey, Oyarzún et al. ( 2019 ) (after noting that they probe larger radii) and Lacerna et al. ( 2020 ). For C and Mg, at low masses the radial gradients are slightly positive, and they become shallower or slightly negative as galaxy mass increases. The gradients are independent of the type of galaxy. For the heavier elements, ETGs display similar gradients at all masses for [Ca/Fe], and become slightly negative as mass increases for [Ti/Fe]. High-mass spirals show larger variation in gradients with mass. The overall trend for Ca is unclear, but for Ti, the gradients become more negative as galaxy mass increases, like ETGs. N gradients go from slightly positive or shallow at low masses, to negative at high masses for both types. Finally, for Na, ETGs display a slight steepening in the gradient for the most massive galaxies while gradients are consistent for both types at intermediate masses. This suggests that the mass or some other property drives Na abundance variation rather than galaxy type. Briefly, we find several interesting features while studying these radial gradients in these parameters for different galaxy types. The most striking difference is in the age gradients. While ETGs display relatively flat age gradients, at all masses, LTGs show negative age gradients, which become steeper with mass. Metallicity behaves exactly the opposite way, with negative gradients steepening with mass for ETGs and negative gradients overall for LTGs. These properties could provide clues into how these different galaxy types and their stellar populations assembled and evolved. C and Mg abundances and radial gradients are remarkably similar for both types, suggesting that star formation time-scales are independent of galaxy type but rather depend on some other physical property i.e. galaxy mass or velocity dispersion. Na behaves in the same way as metallicity, showing negative gradients for LTGs, and negative gradients that steepen with mass for ETGs. We explore relations with velocity dispersion in the following section. 3.2.4 Local versus global relationships Along with radial gradients, it is informative to look at trends with velocity dispersion. We plot the spatially resolved parameters for both galaxy types, in different mass bins, so that we can identify local deviations (radial) from the global relation with velocity dispersion. The increased sample size and inclusion of different morphologies considerably expands our parameter space compared to P19, section 3.3. Fig. 18 shows all our derived stellar population parameters against the velocity dispersion. These velocity dispersion profiles shown in Fig. 5 have been calculated by combining the dispersions from the individual spectra belonging to each bin. The symbols represent three key pieces of information: (i)
Paper: 2019ApJ...872..105S (fulltext)
and discuss the future work needed to improve our understanding of r -process enrichment at high redshifts. 2. Method To model the formation of DNSs, we use published results from the StarTrack (Belczynski et al. 2002 , 2006 , 2008 ) population synthesis code. These simulations are described in detail in Dominik et al. ( 2012 ) and include three major improvements over previous StarTrack versions with regards to stellar winds (Belczynski et al. 2010 ), common envelope (CE) formulation, and compact object formation (Fryer et al. 2012 ). We have analyzed the intermediate data from these simulations that are publicly available. 7 7 The public data from the Synthetic Universe Project are available at: www.syntheticuniverse.org . We briefly summarize the method in this section. Each of the models we analyze has 2 × 10 6 binaries, initialized by four different parameters: (i) the primary star’s mass M 1 ; (ii) the mass ratio q = M 2 / M 1 , with M 2 being the mass of the secondary star; (iii) the semimajor axis a of the binary; and (iv) the eccentricity e . The masses of the primary are drawn from the Kroupa initial mass function (IMF) from 5 to 150 M ⊙ . The mass ratio is assumed to have a flat distribution between q = 0–1 with the minimum mass of the secondary considered to be 3 M ⊙ . The distribution of initial binary separation is assumed to be flat in log( a ), and the eccentricity is drawn from thermal equilibrium distribution Ξ( e ) = 2 e , with e ranging from 0 to 1. 2.1. Model Variations A number of parameters are varied among the models, which are meant to parameterize the uncertainties in our understanding of binary neutron star formation. The first parameter that is varied deals with the behavior of binaries when a star enters into a CE with a Hertzsprung gap (HG) donor star; during a CE, a star enters the envelope of its companion, exchanging orbital energy to unbind the donor’s envelope. For giant stars, with a clear core–envelope boundary, the end result of this process (so long as there is enough orbital energy available to keep the system from merging) is a closely bound binary comprised of the accretor star and the giant star’s core. However, HG stars lack well-defined cores, and studies are inconclusive as to whether binaries entering into a CE during this phase can survive without merging (Deloye Taam 2010 ). We note that allowing HG stars to survive a CE significantly increases the merger rates of double black hole binaries to a level that exceeds current LIGO constraints (Belczynski et al. 2007 ), although it matches NSM rate estimates (Chruslinska et al. 2018 ). Two different submodels, A and B, have been analyzed, and treat CE events with HG donor stars differently. Submodel A treats HG stars such that a core could be distinguished from an envelope in their evolutionary phase, hence a successful CE ejection is possible. Submodel B assumes any system entering into a CE with a HG donor will merge. As a second parameter, the
Paper: 2019A&A...621A...1R (acknowledgments)
The Swedish 1-m Solar Telescope is operated by the Institute for Solar Physics of Stockholm University in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofísica de Canarias. We made use of resources provided by the Swedish National Infrastructure for Computing (SNIC) at the High Performance Computing Centre North at Umeå University. JdlCR is supported by grants from the Swedish Research Council (2015–03994), the Swedish National Space Board (128/15) and the Swedish Civil Contingencies Agency (MSB). This project has received funding from the European Research Council (ERC) under the European Union’s Horizon 2020 research and innovation programme (SUNMAG, grant agreement 759548). This research was supported by the CHROMOBS and CHROMATIC grants of the Knut och Alice Wallenberg foundation. This research has made use of NASA’s Astrophysical Data System.
Paper: 2019A&A...626A.100B (fulltext)
from CO line emission ( Olofsson et al. 2002 ; González Delgado et al. 2003 ), as well as a linear fit to the observed values. These observed wind properties can be compared with wind properties derived from the model grid in order to evaluate what combinations of seed particle abundance and piston amplitude best reproduce the observational data. Constraining the pulsation period in the model data to less than 800 days to fit the periods of the observed sample, we found that models with log( n d ∕ n H ) = − 14.5 and u p = 2.0 give the best agreement with observations. The bottom panel of Fig. 12 shows the distribution of the differences between model mass-loss rates and the expected mass-loss rates from the linear fit to the observed values. The middle panel of Fig. 12 shows mass-loss rate as a function of pulsation period for the same observational sample as shown in the top panel. One way to asses how sensitive the mass-loss rates are to the pulsation mode is to look at the group of stars indicated by blue symbols in Fig. 12 . These stars are semiregular/irregular variables (SRVs), probably pulsating in the first overtone as indicated by their short pulsation periods. Even though these stars presumably pulsate in first overtone, they follow the same correlation between mass-loss rates and wind velocities as the fundamental mode pulsators (red symbols). This is clearly seen in the top panel of Fig. 12 , which shows the mass-loss rates and wind velocities for stars pulsating in different modes. Assuming that our reasoning about the stars with shorter periods being first overtone pulsators is correct, it seems that differences in the pulsation modes may not notably affect the mass-loss rates in M-type AGB stars. 5.2 Trends at the longest pulsation periods Figure 13 shows the mass-loss rate as a function of pulsation period for all models in the grid with a stellar wind. Also shown are observed mass-loss rates estimated from CO-line emission against pulsation periods for M-type AGB stars from Olofsson et al. (2002) and González Delgado et al. (2003) , as well as a sample of OH/IR stars taken from De Beck et al. (2010) and He Chen (2001) , with mass-loss rates estimated from CO-line and dust emission, respectively. In model data and in observations the mass-loss rates level off at longer pulsation periods, even if the models do not reach quite as high as the observed mass-loss rates. However, it is important to remember that the grid is currently biased against models with very high mass-loss rate, where dust formation and wind acceleration is very abrupt, resulting in non-converging dynamical structures (see Sect. 4.1 ). The exact level where the mass-loss rate flattens in the model values might also depend on the inner boundary conditions of the models, for example the value of f L or the phase-shift of the luminosity variation with respect to the radial variation of the surface layers ( Liljegren et al. 2016 , 2017 ). It might also be
Paper: 2020ApJ...894...61P (fulltext)
of LLAGNs, its 2–10 keV intrinsic luminosity is only , or ∼10 −8.5 of the Eddington luminosity of the putative SMBH (Yang et al. 2015 , hereafter Y15 ). Using NSF’s Karl G. Jansky Very Large Array (VLA) observations at a central frequency of 6.6 GHz with a bandwidth of 2 GHz, Y15 detected a weak radio source (with a luminosity of νL ν [6.6 GHz] ) coincident with the optical nucleus of M32, which they argued is most likely the radio counterpart of the SMBH. Together with Sgr A* and M31* (Crane et al. 1992 ; Yang et al. 2017 ), the LLAGN in M32 provides a rare but promising opportunity to study feedback from the most under-luminous SMBHs on parsec-scales. In this work, we present follow-up high-resolution, dual-band VLA observations to further explore the nuclear radio source in M32. Section 2 describes our new VLA observations and data reduction procedures. Section 3 presents the main observational results, in particular, morphological and spectral properties of the nuclear source. Section 4 is devoted to understanding the nature of this source. To be consistent with previous work, we adopt a distance of 780 kpc for M32 (1′′ corresponds to 3.8 pc; Young et al. 2008 ). Quoted errors are at the 1 σ confidence level, unless otherwise stated. 2. Observations and Data Preparation Our VLA A-configuration observations of M32 were initially scheduled for the 2015A semester (Project ID: 15A-049, PI: L. Sjouwerman). However, due to a failure in the antenna control system, only one of the three requested observing epochs was successfully conducted, which resulted in two 3 hr integrations on 2015 July 20, one in C -band and the other in Ku-band. Subsequently, on 2016 December 13, 2016 December 14, and 2017 January 22, we conducted three additional epochs of dual-band observations through a new program (Project ID: 16B-073) with the same setting as in 15A-049. In all four epochs, the C -band (Ku-band) observation was centered at 6.0 (15.0) GHz, with a bandwidth of 4 (6) GHz as captured with the 3 bit samplers. Using a visibility integration time of 2 s, the available bandwidth was divided into 1 MHz dual polarization channels to aid the selective removal of RFI. A phase-referenced mode was adopted, with J0038+4137 serving as phase calibrator every 10 minutes for C -band and every 8 minutes for Ku-band. 3C48 7 7 3C48 underwent a major event since 2016; however, it did not significantly affect our data analysis and result (private communication with NRAO helpdesk). See also https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/pol . was used as flux and bandpass calibrator with a 3% systematic uncertainty (Perley Butler 2013 ). In total, we obtained 8.5 hr on-source integration for C -band and 10.0 hr for Ku-band. A log of the VLA observations is given in Table 1 . Table 1 Log of VLA observations Note. (1) Observation date. (2) Central frequency. (3) Synthesized beam FWHM of major and minor axes, and position angle of the major axis.
Paper: 2019A&A...629A..82T (acknowledgments)
The authors thank the anonymous referee for suggestions that led to important improvements in the paper. The research leading to these results has received funding from the European Union’s Horizon 2020 Programme under the AHEAD project (grant agreement n. 654215). A.T. and E.N.E. would like to thank Boğaziçi University Research Fund Grant Number 13760 for financial support. A.T. acknowledges support by the Study in Italy Grant, Research Under Academic Supervision. F.T. acknowledges support by the Programma per Giovani Ricercatori – anno 2014 “Rita Levi Montalcini”. M.G. is supported by the Lyman Spitzer Jr. Fellowship (Princeton University) and by NASA Chandra GO7-18121X and GO8-19104X. R.S. acknowledges financial contribution from the agreement ASI-INAF n.2017-14-H.0. The authors thank Massimo Cappi for useful discussions. A.T. thanks François Mernier and Jelle de Plaa for useful discussions.
Paper: 2019A&A...628A..65A (acknowledgments)
We are grateful to the anonymous referee for constructive comments that helped us improve the manuscript. We thank Luis Colina, Alvaro Labiano, Leonard Burtscher and Daniel Asmus for useful discussions. A.A.-H. and I.G.-B. acknowledge support from the Spanish Ministry of Science, Innovation and Universities through grant AYA2015-64346-C2-1-P, which was party funded by the FEDER program. A.A.-H. also acknowledges CSIC grant PIE201650E36. A.A.-H., S.G.-B. and A.U. acknowledge support from the Spanish Ministry of Science, Innovation and Universities through grant PGC2018-094671-B-I00, which was party funded by the FEDER program. S.G.-B. and C.R.A. acknowledge the Spanish Plan Nacional de Astronomía y Astrofísica under grant AYA2016-76682-C3-2-P. M.P.-S. and D.R. acknowledge support from STFC through grant ST/N000919/1. T.D.-S. acknowledges support from ALMA-CONICYT project 31130005 and FONDECYT regular project 1151239. S.F.H. is supported by European Research Council Starting Grant ERC-StG-677117 DUST-IN-THE-WIND. C.R.A. also acknowledges the Ramón y Cajal Program of the Spanish Ministry of Science and Technology through project RYC-2014-15779. S.R. and M.V. gratefully acknowledge support from the Independent Research Fund Denmark via grant numbers DFF 4002-00275 and 8021-00130. L.K.H. acknowledges funding from the INAF PRIN-SKA 2017 program 1.05.01.88.04. C.R. acknowledges support from the CONICYT+PAI Convocatoria Nacional subvencion a instalacion en la academia convocatoria año 2017 PAI77170080. D.R. acknowledges the support of the Science and Technology Facilities Council (STFC) through grant ST/P000541/1. AU acknowledges support from the Spanish MINECO grants ESP2015-68964 and AYA2016-79006. This paper makes use of the following ALMA data: ADS/JAO.ALMA.2016.1.00254.S and JAO.ALMA.2016.1.01236.S. ALMA is a partnership of ESO (representing its member states), NSF (USA) and NINS (Japan), together with NRC (Canada), MOST and ASIAA (Taiwan), and KASI (Republic of Korea), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO and NAOJ. Based on observations made with the NASA/ESA Hubble Space Telescope, obtained from the data archive at the Space Telescope Science Institute. STScI is operated by the Association of Universities for Research in Astronomy, Inc. under NASA contract NAS 5-26555.