adsabs/WIESP2022-NER huggingface dataset:
Paper: 2020MNRAS.496.1077P (fulltext)
in Section 5 . To facilitate discussion, we define as a reference value. In order for MI to be disfavoured by less than 3σ, must be lowered by with fixed degrees of freedom. Hence, it would likely require a combination of (1) a different definition of r bar , (2) different values for the mass-to-light ratios, possibly including radial dependence, (3) removal of statistical outliers, and (4) other effects. Regarding point (1): There is an unaccounted bias in that the bins are defined by r bar , a point that is affected by uncertainty. The impact of this bias can be roughly estimated by moving r bar one point towards smaller (larger) radii. Doing so yields a reduction in χ 2 relative to , however with the accompanying degrees of freedom . The reduction in degrees of freedom (galaxies) are due to the criterion of at least two points at small radii. Considering the reduced chi square changes from to after the shift to smaller radii, it can be concluded that although the shift in r bar greatly affects the χ 2 it only slightly reduces the statistical disfavouring of MI. Regarding point (2): The dependence on the mass-to-light ratios may significantly affect g bar and consequently both the expected value of MI and the specification of r bar . The fixed mass-to-light ratios used here ( and ) are not optimal for all galaxies. This is clear from the rotation curve plots in Appendix C , where for some galaxies v bar > v obs for limited radial range (e.g. as is the case for CAMB and NGC 4217). To roughly gauge the effect of radically different mass-to-light ratios, we take (corresponding to galaxies dominated by DM at all radii, Lelli et al. 2017 ), yielding an increase in the χ 2 relative to (the result become slightly more significant). This is, however, just an example and one can imagine changing individual mass-to-light ratios and uncertainties potentially having a great impact on the result. Regarding point (3): Fig. 5 shows the distribution of residuals. Of the galaxies with | R i | > 4 UGC11455, NGC 4217, NGC 3109, ESO563-G021, and D631-7 have quality flag 1 whereas CAMB have quality flag 2 (listed in the SPARC data base). This indicates that the disfavouring of MI is not rooted in bad data (Fig. B1 in Appendix B shows residuals after discarding all galaxies with quality flag larger than 1). That being said, as mentioned in point (2), the fixed mass-to-light ratios used are too large in case of NGC 4217 and CAMB and possibly also affect the innermost radii of UGC11455 and ESO563-G021. It can also be argued that the high residual value in case of ESO563-G021 is at least partially rooted in a peculiar selection of points (e.g. see Fig. C1 in Appendix C ). Hence, although data are of good quality, for several of the significant outliers, there are valid arguments that could at least partially explain the high residuals. To be conservative and test the dependence on the outliers, we proceed to discard all galaxies with | R i | > 4. Doing so yield a
Paper: 2018MNRAS.473.1705S (fulltext)
The x -axis in the line profiles gives the lambda channels and the y -axis is the intensity in arbitrary units. Both [S ii ] lines at 6717 Å and 6731 Å are detected. The decomposition of each profile is indicated by the thin lines. The resulting profile is shown as hollow circles and with numbers. The dotted line represents the sum of all fitted components. We assume the brightest component of the line to be the main velocity component of the SNR. It is marked with a 0 in Fig. 6 . For each emission line, we can also identify two extreme velocity components. We shall refer to these as V min and V max . These are marked with a 2 and a 4, respectively. Table 3 reports the velocity components of the [S ii ] λ6717 Å line. Column 1 is the region ID. Columns 2, 3 and 4 indicate the main and extreme velocity components ( V min and V max ) corrected for the local standard of rest (LSR). Column 5 gives the difference in velocity between the extreme velocity components. Column 6 lists the distance of each region in arcmin to the geometrical centre of the SNR. Column 7 lists the [S ii ]/H α line ratio and finally column 8 indicates the associated filament to which each region belongs. As one can see from Fig. 6 and Table 3 , the radial velocity profiles are, in general, composite, implying large motions as expected in a SNR. They show at least three velocity components. The brightest one corresponds to the systemic motion of the SNR, while two faint velocity components are seen at extreme velocities. The several velocity components found in the velocity profiles of the filaments are probably due to the SNR shock expansion, as well as to contamination due to the interaction of the SNR gas with dense interstellar clouds that could be associated with the H ii regions located to the north and south-west of the remnant. Note the existence in the SNR velocity profiles of a velocity component also identified with those of regions located in the Perseus arm, as we will discuss below. The velocities of the main component vary from −34 to −89 km s −1 . Taking the average of all the velocities of the main components in Table 3 , we obtain a LSR velocity of −59 ± 17 km s −1 as a first estimate of the systemic velocity of the SNR. The averages for the extreme velocity components, in all the boxes where velocity profiles were extracted, are V min ∼ −122 km s −1 and V max ∼ +152 km s −1 . Another way of obtaining the systemic velocity is to average the extreme velocity components ( V min and V max ). Doing this, we obtained a systemic velocity V LSR = −50 ± 6 km s −1 . For some regions, we find an additional velocity component averaging 1 km s −1 . This latter velocity component agrees with the component found by Roberts ( 1972 ) in the theoretical H i profile of the TASS model at 12° pitch angle ( V LSR = 0 km s −1 ) in the direction of SNR CTB 109. Roberts ( 1972 ) attributed this component to gas located very close to the Sun in the Orion spur. Fig. 7 shows a
Paper: 2017ApJ...843...60W (fulltext)
flux ratio . The sum of the model primary and rescaled secondary spectra was then rectified to a unit continuum. We then repeated the CCF and tomography analysis of this model set of spectra in the same way as was done with the observed spectra to derive a relationship between reconstructed secondary CCF peak height and the assumed flux ratio at 1525 Å. Interpolating the observed CCF peak height in this relation led to a monochromatic flux ratio estimate of at 1525 Å. This flux ratio is related to the ratio of stellar radii bywhere is the monochromatic surface flux ratio of the stars and is equal to in the continuum at 1525 Å for the adopted temperatures of the two stars in the TLUSTY models. Using the above relation, the resulting radius ratio is , where the uncertainty includes the range of associated with the error in . Koubský et al. ( 2000 ) estimate that the Be star mass and mean radius of 60 Cyg are and , respectively. If we adopt these values, then the mass, radius, and gravity of the subdwarf are , and , respectively. This estimate of the companion mass places it above the Chandrasekhar limit of , so it is possible that the subdwarf may become a neutron star in the future. If so, then 60 Cyg may be a progenitor of a Be X-ray binary system. The estimated gravity of the subdwarf is higher than the limit in the TLUSTY grid that we used in deriving the best-fit . The relative strengths of the spectral features in the model spectra depend on the ionization balance in the model, so if we had used a higher model, we would probably arrive at a slightly higher estimate of and a somewhat lower value of R 2 . It is remarkable that among the now five known Be+subdwarf systems that 60 Cyg and 59 Cyg are separated by about one degree on the sky. The stars are not known as cluster members, but they may be affiliated with other early-type stars in their vicinity. We formed a list of stars with B spectral classifications from the catalog of Skiff ( 2014 ) that are situated within 2 deg of the midpoint between 59 Cyg and 60 Cyg, and then collected proper motions from the Gaia Collaboration et al. ( 2016 ) and parallaxes from Leeuwen ( 2007 ) for these stars. Table 3 shows those selected with proper motions and parallaxes within of those of 59 Cyg (HD 200120) and 60 Cyg (HD 200310). This collection of six objects may form part of a hitherto unknown co-moving group of equal age stars. If so, then 59 Cyg and 60 Cyg may be at a similar stage of binary evolution because of their common origin and age. Table 3 Co-moving Stars in the Vicinity of 59 Cyg and 60 Cyg Note. The spectral classifications are from the collection by Skiff ( 2014 ), proper motions and are from the Gaia Collaboration et al. ( 2016 ), and parallaxes are from van Leeuwen ( 2007 ). The V magnitude and radial velocity data are from SIMBAD, with the exception of the systemic velocities for 59 Cyg (Peters et al. 2013 ) and 60 Cyg (Koubský et al. 2000 ). We found no evidence of the
Paper: 2016MNRAS.459.4174G (fulltext)
to the analytical Newtonian, adiabatic and steady-state wind model, referred to in the following as NASS, derived by Takahashi Janka ( 1997 ). It provides a simple, fully analytical description of the dynamics of the wind outflow at relatively late times or sufficiently far away from the PNS surface. The NASS wind model relies on the general assumptions listed above, in particular, those of a Newtonian PNS gravitational potential and of an adiabatic expansion. In addition, all the elementary and e − /e + weak interaction processes are assumed to be frozen out, nuclear β-decays not to affect Y e or s , and possible deviations from nuclear equilibrium with regard to strong and electromagnetic interactions to have no influence on the thermodynamical properties of the wind. Under such simplifying assumptions, the NASS model cannot predict any time variations of the entropy s , electron fraction Y e , or of the wind mass-loss rate , which are thus treated as constant input parameters. The basic NASS wind dynamics in the regime under consideration, and in particular for high enough entropies, is well approximated (Takahashi Janka 1997 ; Arnould et al. 2007 ) by where M * is the PNS mass, r the radius, v the velocity, T the temperature. The total energy per unit mass E may be obtained by setting a boundary condition and is usually expressed as a function of the wind energy: E = f w × E wind , where and v s equals the local adiabatic sound speed (Arnould et al. 2007 ). For f w = 1, the solutions correspond to a sonic wind, whereas for f w > 1, they are of sub-sonic wind type and usually referred to as breeze solutions. More details about the wind model can be found in Arnould et al. ( 2007 ). The entropy is dominated by photons, electrons and positrons in the wind so that the so-called radiation entropy is given by where f e = 1 in the high- T limit when electrons and positrons are highly relativistic, and decreases with T for high -values. In this study, it is approximated by , where T 9 is the temperature expressed in 10 9 K (Witti et al. 1994 ). Since the nucleosynthesis in the wind is followed by starting at a temperature of T 9 = 9, the total entropy is essentially equal to the radiation entropy (within a few per cent). Note that the ‘radiation entropy’ is measured here per baryon in units of the Boltzmann constant k B and includes both the photon and lepton contributions, but not the baryonic ones. For the initial conditions adopted here, the latter amounts to typically 9 k B /baryon. In the wind model of Takahashi Janka ( 1997 ), the radial evolution of the velocity, v ( r ), and density, ρ( r ), for specified total energy ( E ) per unit mass of outflow material (or, equivalent, for given f w ), given entropy and given mass-loss rate () is determined by the energy equation (based on Bernoulli's equation, see equation 1 ) and the continuity equation . The temperature then follows from the assumption of constant entropy, which couples
Paper: 2019MNRAS.486.1781R (fulltext)
indicating a harder γ-ray spectrum. The curvature parameter β gives a measure of the presence of cut-off in the spectrum at high energies with a large value of β indicating a sharper cut-off. Therefore, investigation of any changes in the values of α and β parameters during the five epochs can point to changes in the γ-ray spectral shape. A change in the γ-ray spectral parameters during a flaring state could point to a change in the position of the γ-ray-emitting region. The variation of the α and β against flux during the five epochs studied here is shown in Fig. 14 . No trend on the variation of α with flux was seen during the five epochs considered here. However, a plot of β with flux (Fig. 14 ) shows a clear trend of a decrease in the curvature parameter with increasing flux. Linear least-squares fit to the data gives β = (−0.005 ± 0.009) F V + (0.168 ± 0.029) with a correlation coefficient of −0.77. Clearly the value of β is high and low at the lowest and the highest flux levels among the five epochs analysed here. Such a trend is also seen by Ackermann et al. ( 2015 ) on analysis of the FSRQs in the third catalogue of AGNs detected by Fermi (3LAC). The changes in the γ-ray spectral shape during different brightness states of 3C 454.3 could point to the emission site located at different regions during different brightness states; however, other explanations could not be ruled out. Studies on the flares of 3C 454.3 in December 2009 and November 2010 that fall in the epochs A and B studied here concluded that the γ-ray emission regions were located close to the central black hole (Jorstad et al. 2012 ; Isler et al. 2013 ). Figure 14. Variations of the parameters α and β with flux. 4.2 Connection between optical and GeV flux variations In a majority of the multiwavelength monitoring observations of blazars, close correlations between the flux variations in different bands were noticed (Bonning et al. 2009 ). This was explained on the co-spatiality of the emission regions emitting in different bands and the correlated optical and GeV flux variations can be understood in the standard leptonic emission processes according to which the same relativistic electrons produce optical and γ-ray emission via synchrotron and IC processes. However, there are a handful of blazars where the emission in the optical and GeV γ-ray bands are found not to be correlated (Chatterjee et al. 2013a ; Dutka et al. 2013 ; Cohen et al. 2014 ; MacDonald et al. 2015 ). During the 9 yr of monitoring data analysed here we found four flaring epochs in the optical, namely A,B, D, and E. During epochs A and B, the optical flare is accompanied by a γ-flare, while at the other two epochs D and E, though the optical flares have amplitudes similar to those of epochs A and B, the γ-rays during epochs D and E were either weak or undetected. This is clearly seen in Fig. 15 , where the logarithm of γ-ray flux is plotted against the logarithm of optical flux. The results of the linear
Paper: 2016MNRAS.462.1847B (fulltext)
a radius of 3 arcmin. The Galactic north/south line runs through azimuthal angles 35°/215°. Of the profiles in Fig. 8 , the maximum from ring 2 appears to align the closest with the Galactic north direction, as might be expected if the column density were increasing towards the Galactic plane. Further discussion is deferred to Section 4.1 . VP15 also noted the presence of azimuthal structure in the XRT halo data; however, they did not separate out the behaviour of rings 1 + 2 or 3 + 4, or proceed to discuss the findings in any depth. 2.3 Halo spectral analysis Spectra from the rings were extracted from the longest PC mode snapshots, or adjacent snapshots, using annular extraction regions tuned specifically for each case. Given their relative separations, spectra from rings 1 and 2 were extracted as one, which increased the statistical quality of the data. The same was done for rings 3 and 4. PC event grades 0–12 were selected for the spectra. As the rings dominate the field of view and leave limited space on the detector from which to extract a good quality background spectrum, the PC mode observations from 2012 (total exposure 132 ks) were used to obtain an accurate non-halo background spectrum, which was extracted using an annulus of size 2.95–7.66 arcmin centred on the source, after five faint background point sources were removed. (Note, the background sources were too faint to have any effect on the radial profiles in Section 2.1 .) Also, as the central source was often observed offset from the centre of the CCD, a given ring samples different parts of the telescope vignetting. To correct for this in the ancillary response files (ARFs), we followed the procedure outlined by Moretti et al. ( 2011 ). This involved using a vignetting map, created by dividing exposure maps made with and without the vignetting correction (calculated at 1.5 keV) applied, to work out the mean vignetting correction in each extraction annulus. Then, using the analytical description of the vignetting given by Tagliaferri et al. ( 2004 ), these were converted to mean off-axis angle estimates that were used as an input to the XRT software task xrtmkarf to apply the energy-dependent mean vignetting correction to the on-axis ARF. The non-vignetted exposure map was also used to estimate the mean exposure in the annulus and the spectral file exposure time was updated accordingly. The spectra were fit in xspec (version 12.8.2; Dorman Arnaud 2001 ), using its implementation of the Cash ( 1979 ) statistic (C-stat), which, when a background spectrum is supplied, is a profile likelihood statistic. The spectra for a given ring at different epochs were modelled using an absorbed power law, with the photon indices and column densities tied between spectra. The absorption was provided by the tbnew model, 3 using the abundances of Wilms, Allen McCray ( 2000 ) and cross-sections of Verner et al. ( 1996 ). The fitting was performed over the energy range 1–5 keV, where the low
Paper: 2020MNRAS.494.4695Q (fulltext)
of Fig. 5 , since the contact between the QS and HS surfaces does not occur yet, QS and HS surfaces coexist independently and we use ‘QSHS’ to denote this stage. When = −1.502164, the contact between the QS and HS surfaces begins to occur, and we define this critical value as the upper boundary of of the TS motion, . In the blue region of Fig. 5 , the compound isosurface is made up of QS, TS, and HS portions, and we use ‘QS + TS + HS’ to denote this stage. Figure 5. Evolution of isosurfaces with i m = 15°, where blue points denote values used in the following figures. Fig. 6 displays phase structures with (a) i m = 45° and (b) i m = 75° in the torus space, where the colours denote different values of . Compared with the result in Fig. 2(b) , we can observe that i m has a significant influence on the phase structure. The evolution of isosurfaces of i m = 45° and 75° with the decrease of are shown in Fig. 7 . Different the result in Fig. 5 , we can observe that when i m = 45°, is larger than . Numerical computation indicates that when i m = 29.65°, is equal to ; if i m > 29.65°, is larger than . QS and TP surfaces coexist independently in the yellow region of Fig. 7(a) , and when the contact between the QS and TP surfaces begins to emerge, the isosurface is made up of QS, TS, and TP portions (see the blue region in Fig. 7a ). If further deceases, the TP portion on the compound surface will be replaced by the HS portion. Numerical computation also shows that with the increase of i m , gradually decreases and approach , so the yellow region in Fig. 7(a) shrinks with the increase of i m . When i m ≥ 61°, is equal to and independent TP isosurfaces totally disappear. Hence, as we can see from the result of i m = 75° in Fig. 7(b) , the compound isosurface in the blue region is made up of QS and TS portions. Figure 6. isosurface with (a) i m = 45° and (b) i m = 75° in the torus space, where the colours denote different values of . Figure 7. Evolution of isosurfaces with (a) i m = 45° and (b) i m = 75°, where blue points denote values used in the following figures. Fig. 8 displays collision curves of i m = 15°, 45°, and 75° in (a) torus space and (b) ω − e projection. From Fig. 8(a) , we find that the φ range of collision curves starting from the point P = ( e 0 , 0, 0) expands with the increase of i m . In the ω − e projection, we can observe that with the increase of i m , the angle between the branches of collision curves increases. In the following section, we will demonstrate that collision curves play an important role in the co-orbital motion and the above findings can explain some co-orbital phenomena. Figure 8. Collision curves of i m = 15°, 45°, and 75° in (a) torus space and (b) ω − e projection. 4 LOCI IN THE TORUS SPACE In this section, we implement numerical integration in the Sun–Jupiter CRTBP to demonstrate the validity of the semi-analytical analysis in the last section and the detail of the co-orbital motion. Numerical integration is
Paper: 2018ApJ...853..188A (fulltext)
et al. 2016a ) in the standard canonical inflationary scenario ( c s = 1), as has been explicitly shown in Rezazadeh et al. ( 2015 , 2016 , 2017a ). In the present paper, we aim to refine the power-law, intermediate, and logamediate inflationary models in light of the current observational results by considering them in the DBI scalar field setting. Motivated by Spaliński ( 2008 ) and Tsujikawa et al. ( 2013 ), we assume in our work that the sound speed is constant during inflation. We examine the viability of each scale factor in light of the Planck 2015 results (Ade et al. 2016a ) and also present some observational constraints on them. The outline of this paper is as follows. In Section 2 , we review briefly the basic equations governing the background cosmology in the DBI inflation. In addition, we also represent the power spectra of the primordial scalar and tensor perturbations. Then, in Sections 3 – 5 , we respectively investigate the power-law, intermediate, and logamediate DBI inflationary models with constant sound speed. Finally, in Section 6 , we summarize our conclusions. 2. DBI Inflation Within the framework of Einstein’s gravity, the action of non-canonical scalar field models is given by (Armendáriz-Picón et al. 1999 ; Garriga Mukhanov 1999 )where g and R are the determinant of the background metric and the Ricci scalar, respectively. Throughout this paper, we work in units where the Planck reduced mass is taken equal to unity, = 1, for the sake of simplicity. In the present work, we concentrate on the DBI non-canonical scalar field with the Lagrangian density (Alishahiha et al. 2004 ; Silverstein Tong 2004 )where is the warp factor, while , and are the canonical kinetic term and potential of the DBI scalar field φ , respectively. For the DBI Lagrangian ( 2 ), the energy density and pressure of the non-canonical scalar field φ are given by where the subscript, “,” indicates the partial derivative with respect to X . Moreover,is the parameter determining the relativistic limit of brane motion in a warped background. Note that the parameter γ is defined in analogy of the Lorentz factor appearing in relativistic particle dynamics. The sound speed is generally defined bywhich specifies the propagation speed of the inflaton fluctuations relative to the homogeneous background. In order for the sound speed to be physical, it should be real and subluminal, (Franche et al. 2010 ). Applying Equations ( 3 )–( 5 ), the sound speed in DBI inflation takes the form For the homogeneous and isotropic background spacetime, we consider the spatially flat Friedmann–Robertson–Walker (FRW) metric. Therefore, varying the action ( 1 ) with respect to the metric yields the Friedmann equations in the DBI inflation as follows: Here, is the Hubble parameter, and a is the scale factor a of the universe. The dot denotes the derivative with respect to the cosmic time t . Considering the FRW metric, the canonical kinetic term is simplified to . Taking the