adsabs/WIESP2022-NER huggingface dataset:
Paper: 2018A&A...620A.183J (fulltext)
of the proxy of the chromospheric heating and maxima of green lines show the regions with the maximum of the proxy of the photospheric heating. The arrows connecting these maxima show their spatial displacement. In three cases out of four, the arrow orientation is in agreement with our estimation of the spatial displacement of the photospheric and chromospheric heating. The spatial separation of the maxima is between 0ʺ̣5 and 1′′. We note that the local maxima of the black lines for the orange and green streamlines located around 3′′ and 2ʺ̣5, respectively, are caused by the evolution of the light bridge intensity and are not related to the WLF. Only in one case (cyan streamline) does the arrow point in a direction that is not in agreement with our interpretation of the projection effect. It should be noted that although the results presented in Fig. 12 are mostly in agreement with our expectations, they cannot be taken as confirmation of the projection effect. The used proxies do not represent the chromospheric and photospheric heating accurately. The magnetic field configuration is more complex and not retrieved precisely with our height-independent model atmosphere. The energy deposition along the field lines is dependent on the physical properties of the chromosphere that are unknown and can differ from pixel to pixel. In addition, the location of the flare-loop footpoints where beams of accelerated particles propagated is not known because imaging hard X-ray data are not available during the time of the Hinode flare scan. Thus, the nature of the atmospheric heating cannot be identified either. 5. Conclusions We use a unique dataset obtained by the Hinode satellite capturing the X9.3-class WLF ribbon with a spectropolarimetric raster scan. The observed Stokes profiles allow us to interpret the temperature stratification of the solar photosphere during this WLF under the assumption of LTE conditions. Comparing these data with temperatures obtained in the post-WLF phase, we can also estimate the heating of the photosphere during the WLF. In the studied WLF ribbon, the flare affected also the upper photosphere. In pixels with the most enhanced emissions, the photosphere is heated down to log τ = −0.5, i.e. approximately 80 km above the solar surface. This is in agreement with the shape of the Stokes I profiles in these pixels as the Fe I lines show absorption features in the far wings, i.e. it signifies a temperature decrease around log τ = 0. Nevertheless, we observe enhancements of the continuum intensity equivalent to that can be ascribed to the WLF ribbon. This corresponds in certain pixels to an increase of I c by 200% (see Figs. 3 and 9 ). Following the analogy of the study by Heinzel et al. (2017) , we assume that this continuum contribution is mostly due to the Paschen recombination continuum formed in the chromosphere. However, we are aware that there are studies locating the WLF sources at various layers in the photosphere (
Paper: 2016MNRAS.463.2425M (fulltext)
Θ s ⊥ ( r ) are denoted as and . The pseudo s are linear combinations of the underlying s. The linear transform relating the two power spectra e.g. and , or, equivalently and , are encapsulated in a mode-mixing matrix. However, due to our assumption of factorizability of the mask we can write the mixing matrix as product of two different mode-mixing matrices. The radial and angular mode-mixing matrices denoted as W r and W a depends on the w r and w a , respectively and can be expressed as follows: For nearly all-sky coverage the mixing matrices can be inverted to recover the underlying s. However, binning may be required for partial sky coverage. The power spectra and can be dealt with in an exactly similar manner. The mixing matrices remain unchanged. It is possible to further simplify the expressions by using an approximate form for j l ( x ) valid at high l (LoVerde Afshordi 2008 ): A few comments are in order. The PCL based sub-optimal estimators for projected or 2D data sets were developed by Hivon et al. ( 2002 ) in the context of CMB power-spectrum estimation. It has also been generalized for computation of cross-spectra involving external data sets (Munshi et al. 2011a ). More recently, it has been extended for the 3D data by (Munshi et al. 2011b ) and was successfully used for analysing 3D weak lensing power spectrum using CFHTLenS 19 data (Kitching et al. 2014 ). Though PCL based estimators remain sub-optimal, suitable use of weights can improve their performance. The results presented here are generalization of the one presented in Munshi et al. ( 2011b ) (where the effect of radial selection function was not included; i.e. W r = 1.) and will have general applicability in cross-correlating arbitrary projected field (3D) against a 3D cosmological tracer. Similar results can be derived for higher-order spin fields by suitably modifying . Generalization to arbitrary spin. For the sake of completeness we also provide the expression for a generic 3D field with an arbitrary spin S α. We would like to cross-correlate it with another arbitrary projected field with spin S ′ say β. In this case the basis of expansion involves spinorial harmonics . The spin-harmonics are defined in terms of D matrices as (Penrose Rindler 1984, 1986 ; Varshalovich, Moskalev Khersonskii 1988 ). The forward and backward transformations are obtained by replacing in equations ( 11 ) and ( 12 ) with spin-weighted spherical harmonics or spinorial harmonics : We follow the same notation as before but we assume the 3D field is a spinorial field of spin-weight S and is a projected filed of spin S ′ . The harmonic components of these fields are now denoted as S α l m and . The suffix s denotes quantities defined in partial sky. The radial component of the mixing matrix given in equation ( A4 ) remains unchanged. However, the angular mode-mixing matrix depends on the spin of the specific field equation ( A5 ). The corresponding auto- and cross-spectrum are defined as
Paper: 2016A&A...589A.113A (fulltext)
mass vs. period of h Per members, including detected and non-detected X-ray sources. For MS stars the separation between saturated and non-saturated regimes occurs at fixed values, regardless of stellar mass. When is computed as , this threshold occurs at ( Wright et al. 2011 ), which turns into when it is taken into account that . Therefore, assuming that the same applies to intermediate-age PMS stars and considering the internal structures of 13 Myr old stars of different masses, it is possible to compute rotational periods corresponding to this value. We report in Fig. 7 this expected separation between non-saturated and saturated regimes (blue dashed line). We also computed where the supersaturated regime is expected to occur. However, the paucity of observational constraints of supersaturation makes this prediction more uncertain. Hence we considered the two physical mechanisms suggested to explain the supersaturation: centrifugal stripping ( Jardine Unruh 1999 ) and polar updraft migration ( Stȩpień et al. 2001 ). The centrifugal stripping mechanism predicts that coronal emission might be reduced when rotational velocity increases because the largest coronal structures might be disrupted by the centrifugal force. This mechanism takes place when the largest coronal structures extend beyond the co-rotation radius 4 . Assuming therefore that centrifugal stripping becomes effective when the co-rotation radius is smaller than (i.e., assuming that coronal structures extending up to above the photosphere), as suggested by Wright et al. (2011) , we obtained the dash-dotted red line shown in Fig. 7 . The polar updraft effect instead predicts that when the rotational velocity increases, then the magnetic flux emergence becomes more efficient near stellar poles, leaving the equatorial regions free of magnetic flux tubes and hence without coronal structures. Wright et al. (2011) quantified the efficiency of this mechanism by defining the parameter where and are the mass and radius of the radiative core, is its rotational angular velocity, and is the colatitude on stellar surface. Considering MS stars, they suggested that supersaturation occurs where . Considering again the internal structures of 13 Myr old stars, assuming that the core rotates with the same period as the convective envelope and averaging over the stellar surface, we computed the locus corresponding to and report this threshold with a red dashed line in Fig. 7 . A more detailed discussion on centrifugal stripping and poleward migration is reported in Sect. 5 . Fig. 7 Mass vs. period of X-ray selected h Per members detected (black) and undetected (gray) in X-rays. Filled and open symbols indicate single and binary stars, respectively. Red lines separate the loci corresponding to supersaturation and saturation, assuming centrifugal stripping (dash-dotted line) or polar updraft (dashed line) theories. The blue dashed line marks the expected separation between non-saturation and
Paper: 2016MNRAS.460.2675R (acknowledgments)
We are grateful to an anonymous referee for helpful comments and suggestions that improved this work. JR thanks Lister Staveley-Smith for useful comments. We thank Claudia del P. Lagos and Han-Seek Kim for providing their model predictions used in Fig. 11 . We also thank the staff of the GMRT for their assistance. The GMRT is operated by the National Centre for Radio Astrophysics of the Tata Institute of Fundamental Research. This research was funded by an Australian Indian Strategic Research Fund (AISRF) grant. This fund was jointly administered by the Department of Innovation, Industry, Science and Research in Australia and by the Department of Science and Technology in India. The project title was ‘Gas in Galaxies in the Distant Past’. Parts of this research were conducted by the Australian Research Council Centre of Excellence for All-sky Astrophysics (CAASTRO), through project number CE110001020. The G10/COSMOS cutout tool uses data acquired as part of the Cosmic Evolution Survey (COSMOS) project and spectra from observations made with ESO Telescopes at the La Silla or Paranal Observatories under programme ID 175.A-0839. The G10/COSMOS cutout tool is hosted and maintained by funding from the International Centre for Radio Astronomy Research (ICRAR) at the University of Western Australia. Full details of the catalogue can be found in (Davies et al. 2015 ) or on the G10 website at http://ict.icrar.org/cutout/G10/ .
Paper: 2017A&A...608A..18G (fulltext)
to distance determinations made with the Baade-Wesselink method, and bias the calibration of the IR period–luminosity relation. Our previous works ( Mérand et al. 2006 ; Gallenne et al. 2011 , 2013 ) showed that these CSEs have an angular size of a few stellar radii and a flux contribution to the photosphere ranging from 2% to 30%. While in the near-IR the CSE emission might be negligible compared with the stellar continuum, this is not the case in the mid- and far-IR, where the CSE emission dominates. Mérand et al. (2007) and Gallenne et al. (2011 , 2013) pointed out a possible correlation between the pulsation period and the CSE brightness in the near- and mid-IR. It seems that long-period Cepheids have brighter CSE than shorter period Cepheids. This could indicate a mass-loss mechanism linked to the stellar pulsation, and therefore that heavier stars experience higher mass-loss rate. This also seems to be consistent with the theoretical work of Neilson Lester (2008) who found a correlation between the mass-loss rate and the pulsation period. As discussed in Sect. 3 , our SPIPS fits from to band provide estimates of IR excess at m and m for Cepheids having Spitzer light curves ( Scowcroft et al. 2011 ). The derived excess magnitudes are listed in Table 2 . Examples of our final fit are shown in Figs. A.1 , A.2 , and A.4 . We clearly see an offset above the photospheric level (in grey), and the stars appear brighter due to the presence of IR excess. In Fig. 7 we plotted those IR excess with respect to the pulsation period, and we note that all stars exhibit an excess. However, the trend seems constant with the period at these wavelengths, which differs from our previous studies at 2.2 and m for Galactic Cepheids. We performed a linear and constant fit for both wavelengths combining all Magellanic Cloud Cepheids. From an -test, we found that a constant model is more statistically significant. We estimated an average excess of mag and mag (again, we took the standard deviation). These excesses are particularly significant and might be the signature of circumstellar dust. These results are relevant for the mid-IR calibration of the Cepheid distance scale (see e.g. Freedman et al. 2012 ). Some authors find that there are several advantages to using a mid-infrared relation. The effect of the dust extinction is reduced, the intrinsic dispersion is smaller than in the band, and the metallicity effect is supposed to be minimal. However, going to longer wavelengths brings new systematics that impact the distance scale, and which are usually not taken into account. Our measured IR excesses convert to a systematic error of 2.6% and 3% on the distance at 3.6 and m, respectively. Unfortunately, recent papers ( Scowcroft et al. 2016 ; Freedman et al. 2012 ; Monson et al. 2012 ) do not take into account this additional uncertainty. On the other hand, their relation is calibrated for a Cepheid plus a CSE, which is still valid if we assume that all Cepheids
Paper: 2021ApJ...913...18B (fulltext)
14.0, is notable. If O vi arises primarily at interfaces between cool clouds and a hot, volume-filling phase, the lack of O vi absorption suggests partial covering of such structures in the relevant galaxy halo(s) and intragroup medium. For the stronger of the two H 2 -bearing velocity components, the excitation state of the molecules may be due to UV pumping by an ambient radiation field that is elevated above that observed in the solar neighborhood by a factor of α = 10–100. This field may be produced by shocks arising from collisions between cool clouds or at the interface between these clouds and the hot, ambient medium, or may imply the presence of undetected star formation obscured by dust or the glare of the QSO (see point III below). In this picture, the observed rise in b value with J level can be attributed to stratification in the cloud, in which J > 1 levels are primarily populated closer to the surface of the cloud and thus in more diffuse and spatially extended gas than the lower rotational levels (e.g., Abgrall et al. 1992 ; Balashev et al. 2009 ). The galaxy redshift survey presents an unprecedented opportunity to explore the origin and physical conditions of an H 2 -bearing DLA in the context of its galactic environment. Nine galaxies are detected within d = 600 pkpc and Δ ν g = 300 km s −1 of the absorber, and here we discuss the possible association of the diffuse molecular gas and the relevant galaxies. I. Association with a massive, early-type galaxy at d = 41 pkpc . Several lines of evidence favor an association between the H 2 -bearing DLA and the massive elliptical galaxy, G41. In addition to having the smallest projected distance ( d = 41 pkpc) and a modest velocity offset (Δ ν g = −57 km s −1 ), this galaxy is also the most massive system observed at the redshift of the DLA and represents almost 70% of the stellar mass of the small group comprising this galaxy and its four closest neighbors. Furthermore, as discussed in Section 4.3 , the chemical abundance pattern suggests Fe enhancement consistent with significant contributions from Type Ia SNe. This chemical signature is commonly seen in the inner halos of massive quiescent galaxies ( d ≲ 100 pkpc), in contrast to their outer halos and star-forming systems (e.g., Zahedy et al. 2017 ). The observed subsolar metallicity at d ≈ 40 pkpc is not in tension with the supersolar stellar metallicities often observed at the centers of massive quiescent galaxies, due to the stellar metallicity gradients observed in these systems (e.g., Spolaor et al. 2010 ) and the significant scatter seen within single luminous red galaxy halos suggesting a range of gas cloud origins and poor gas mixing (Zahedy et al. 2019 ). Given the observed projected distance, d = 41 pkpc, the absorber could be associated with either the outer ISM or true CGM of this massive galaxy. Significant amounts of neutral atomic and molecular gas have been observed in early-type galaxies at low and intermediate
Paper: 2018A&A...614A..23K (fulltext)
~ 2–3 × 10 12 G) magnetic field strength of SMC X-3 and for cm ( Ghosh et al. 1977 ; Ghosh Lamb 1979a , b , 1992 ), where α is a constant that depends on the geometry of the accretion flow, with α = 0.5 the commonly used value for disk accretion, Ṁ 15 is the mass accretion rate in units of 10 15 g s −1 (estimated for the observed luminosity of ~ 1.5 × 10 38 erg s −1 and assuming an efficiency of 0.2), B 9 is the NS magnetic field in units of 10 9 G, M 1.4 is the NS mass in units of 1.4 times the solar mass, and R 6 is the NS radius in units of 10 6 cm. Three prominent broad emission lines are featured on top of the broadband continuum. Centered at ~ 0.51 keV, ~ 0.97 keV, and ~ 6.64 keV, the lines are consistent with emission from hot ionized plasma. The 6.6 keV line is most likely due to K-shell fluorescence from highly ionized iron, while the 0.5 keV line is most likely the K α line of the N VII – Ly α (500) ion. The 1 keV line is most likely the result of combined Ne K α and Fe L-shell emission lines. Both the 1 keV and 6.6 keV lines appear to be broadened, although as we discuss below, the broadening is most likely artificial and is the result of blending of unresolved emission lines from relevant atoms at different ionization states. Both the 1 kev and the 6.6 keV lines are often detected in the spectra of XRBs and appear to be particularly pronounced in the spectra of X-ray pulsars (e.g., Boirin Parmar 2003 ; Díaz Trigo et al. 2006 ; Ng et al. 2010 ; Kolehmainen et al. 2014 ). While in many XRBs the emission lines (and particularly the iron K α line) are attributed to reflection of the primary emission from the optically thick accretion disk (e.g., Gilfanov 2010 , and references therein), in the case of X-ray pulsars (and in this source), the emission line origin is more consistent with optically thin emission from rarefied hot plasma that is trapped in the Alfvén shell of the highly magnetized NS (e.g., Basko 1978 ). The (apparent) line broadening is consistent with microscopic processes, i.e., Compton scattering for a scattering optical depth of 0.5–1 (e.g., Eq. (30) in Basko 1978 ), rather than macroscopic motions, such as the rotation of the accretion disk or the surface of the donor star. The measured width of the two lines corresponds to line-of-sight velocities of ~ 1.6 × 10 4 km s −1 . For the disk inclination of ≲45°, expected for SMC X-3 ( Townsend et al. 2017 ), this corresponds to a 3D velocity of ≳ 2.3 × 10 4 km s −1 . Assuming Keplerian rotation, these velocities would correspond to a distance of ≲ 350 km from the NS, which is almost an order of magnitude smaller than the estimated magnetospheric radius. Furthermore, the phase-resolved pn spectra and the RGS spectra (discussed below) indicate that the line broadening may be an artifact, resulting from the blending of unresolved thin emission lines at different centroid energies. 4.3 Short-term evolution of the spectral continuum and emission lines The timing analysis of
Paper: 2018MNRAS.479.2308B (acknowledgments)
ACKNOWLEDGEMENTS We are grateful to the anonymous Referee for his/her careful reading of the manuscript and a number of comments, which improved the quality of this manuscript. Based on data products from observations made with ESO Telescopes at the La Silla Paranal Observatory under ESO programme ID 179.A-2005 and on data products produced by terapix and the Cambridge Astronomy Survey Unit on behalf of the UltraVISTA consortium. Based on data obtained with the European Southern Observatory Very Large Telescope, Paranal, Chile, under Large Programs 175.A-0839 (zCOSMOS), 179.A-2005 (UltraVista), and 185.A-0791 (VUDS). This work was partially supported by the Consejo Nacional de Investigaciones Científicas y Técnicas (CONICET) and the Secretaría de Ciencia y Tecnología de la Universidad de Córdoba (SeCyT).
Paper: 2016MNRAS.462.4157A (fulltext)
amount of condensation. It appears that the introduction of thermal conduction has been helpful to restrain the numerical diffusion. As we said in Section 4.1 , thermal conduction is a diffusive process: like numerical diffusion, it destroys the smaller structures smoothing the temperature gradients and creating a homogeneous mixture at intermediate temperature. However, unlike numerical diffusion, the scale at which thermal conduction operates does not depend on the grid resolution, but on physical conditions of the problem, in particular on the temperature of the hot coronal gas (see equation 1 ). Thus, although at very low resolutions numerical diffusion is the dominant diffusive process, increasing the resolution, thermal conduction can become the dominant process. The development of turbulence in the wake of our fountain clouds produces a hierarchy of smaller and smaller structures, following a power-law distribution. The size of the smallest structures is determined by resolution, however their survival is largely determined by thermal conduction which operates efficiently on scales below the Field length (∼ 20 pc). Therefore, the smallest clouds (a few parsec) do not influence the growth of cold gas mass because they are not able to survive, mix with hotter gas and trigger condensation. The convergence in condensation shown in Fig. 4 may reflect this change of regime. After these considerations we can conclude that the resolution 2pcx2pc is good enough for our kind of simulations and we proceed with this resolution for the rest of the paper. 4.4 Thermal conduction effects at different <inline-formula><tex-math id="TM0019" notation="LaTeX"></tex-math></inline-formula> Thermal conduction strongly depends on the temperature of the hot gas. The heat conduction flux, , is proportional to T 5/2 (see equation 2 ), then small differences of temperature lead to large differences of exchanged heat. The consequence is that at high temperatures cold structures can evaporate much faster. However, as we saw in Section 2 , when the temperature changes on scales smaller than the mean free path of the conducting electrons, thermal conduction enters in a regime of saturation and the correlation between high temperatures and high heat transfer can be mitigated. A criterion to ascertain the effects of saturated thermal conduction on a structure with length scale l is given by the global saturation parameter: For σ 0 ≳ 1, the effect of saturated thermal conduction becomes significant (Cowie McKee 1977 ; Dalton Balbus 1993 ). In our simulations, the effect of saturation is negligible for the global evolution of the cloud: σ 0 ≪ 1 for l = R cl = 100 pc, regardless of coronal temperature. However, the saturation effect on the cold cloudlets in the turbulent wake could become more and more important (σ 0 ∝ l −1 ) for increasing coronal temperature. Only at T cor = 8 × 10 6 K σ 0 ≲ 1 and at these high temperatures the saturation effect could have a slight