adsabs/WIESP2022-NER huggingface dataset:
Paper: 2021MNRAS.503.5692W (fulltext)
experiments in Rongcheng, along with the measured median profiles, are shown in Fig. 4 . The performance for estimation by the Dewan outer scale is the worst among four models. The estimated by the Dewan model is much smaller than the measured value below 8 km, with an abrupt change occurring at 8 km, and much larger than the measured value above 8 km. For completeness, Appendix A discusses the individual profiles in Rongcheng. The HMNSP99 and Thorpe models have a similar performance for estimation below 12 km, and are consistent with the measured values in both change trend and magnitude. At high altitude, the two models deviate from the measured values by different degrees: a larger deviation above 12 km appeared for sounding experiments for the Thorpe model. The Sθ model has the best prediction results, and it is consistent with the observed median profile in variation trend and magnitude. The relative error profiles for the above four models are illustrated in Fig. 4(b) . The relative error can demonstrate the variation of the deviation between the estimated value and the measured value with altitude. The deviation between the Dewan model and the observed median turbulence is the largest, while that of the Sθ model is the smallest, for all altitudes. The HMNSP99 and Thorpe models have a similar performance in RE (relative error) below 10 km, while HMNSP99 has better results above 10 km. Figure 4. Evaluation of the estimated median profiles (Rongcheng): (a) comparison of the estimated median profiles using the Thorpe (red), HMNSP99 (blue), Dewan (purple) and Sθ (green) models, and the measured profiles (black); (b) relative error of the estimated median profiles. Table 4 presents the RMSE, R xy , and MRE for the estimations of median obtained from the considered four outer-scale models for balloon experiments. The RMSE and MRE of estimated by the Sθ model are better than those for the other three models. It is worth noting that the R xy values of Thorpe, HMNSP99 and Sθ exceed 0.9, which means that the three models may accurately reflect the variation trend of measured turbulence. For the median values, the correlation coefficient of the estimated by the Sθ scale was the highest of all four models. As shown in Table 4 , the average relative error of the median values was generally lower than 1.6%, which meets estimation accuracy requirements for . Table 4. Statistical analysis for the estimation of in Rongcheng using the four outer-scale models, given for median profile. The specific results of the integral statistics of each sounding balloon and median profiles for Rongcheng are tabulated in Table 5 . Clearly, the Sθ scale provides reliable results, close to observed values. As can be seen, our new outer-scale model gives reasonable results when calculating integrated atmospheric parameters. Table 5. Comparison of integrated atmospheric parameters obtained via values observed by meteorological balloons and estimated by the Sθ model
Paper: 2020ApJ...894....8P (fulltext)
1843.03. Contours are spaced in 5% intervals in core mass fraction, with darker colors corresponding to high core mass fraction and lighter colors corresponding to low core mass fraction. As expected, denser planets with higher core mass fractions survive to shorter orbital periods. The limiting orbital period of a pure iron planet is approximately 3.8 hr at , 3.6 hr for , and 3.5 hr for . We note that this figure marginalizes over stellar mass because stellar mass only weakly affects the Roche-limiting minimum orbital period (Rappaport et al. 2013 ). The stellar density does, however, affect whether the planet can reach its Roche limit before colliding with the star (i.e., a = R ⋆ ) and thereby the minimum survivable orbital period for the planet. Figure 4. Composition and mass constraints on planets with KOI 1843.03's measured transit radius and host star mass. Colored contours show the values of the planet mass, which increases with decreasing orbital period and increasing core mass fraction. The boundary of the colored contours in the lower left corner corresponds to the Roche limit. The dashed gray line indicates the measured orbital period of KOI 1843.03. Figure 4 displays interpolated planet masses for a range of orbital periods and core mass fractions in systems consistent with KOI 1843.03's host star mass and transit radius. The boundary of the colored contours in the lower left corner corresponds to the Roche limit; as orbital period decreases, the core mass fraction is more tightly constrained. As anticipated, considering a fixed orbital period in Figure 4 , increasing the planet’s iron mass fraction increases the planet’s mass. Less intuitively, at fixed core mass fraction, decreasing the planet’s orbital period also leads to an increase in the inferred planet mass. This is due to the tidal distortion of the planet; at shorter orbital periods, the volume of the planet exceeds by larger and larger factors. As it orbits less than one stellar radius from the host star’s photosphere ( a / R ⋆ = 1.9), KOI 1843.03 will be significantly elongated in the direction of the star due to tidal distortion. Based on the parameters of KOI 1843.03, our models predict that it must be at least 31% longer along the star–planet line than along the perpendicular axes (aspect ratio of about 1.3), and our models support a value up to nearly twice as long along the star–planet line (aspect ratio of almost 1.8); various possibilities are illustrated in Figure 5 . For comparison, Saturn has an aspect ratio of about 1.1. Figure 5. Aspect ratio constraints on KOI 1843.03, as a function of core mass fraction. To the left of the plot, we indicate the maximum and minimum values of aspect ratio supported by the model grid and show three representative planet cross-sections (assuming ellipsoidal shapes); peach and navy represent the iron-dominated core and silicate mantle, respectively. Results are shown both at the measured transit radius 0.61 and the 1 σ limits
Paper: 2020ApJ...895...61R (fulltext)
structure is completely determined by the conservation of rest-mass, energy, and angular momentum of the accreting gas, the Polish donut model is less constrained. It requires the relation between the fluid angular velocity and the fluid angular momentum per unit energy, Ω = Ω ( l ), as well as the fluid angular momentum per unit energy at the inner edge of the disk, l in . The simplest Polish donut model is the configuration with l = constant, which turns out to be marginally stable with respect to axisymmetric perturbations (Font Daigne 2002 ). Accretion disks in the Kerr spacetime only exist for l ms l l mb , where l ms and l mb are, respectively, the angular momentum per unit energy of the marginally stable equatorial circular orbit (or ISCO) and of the marginally bound equatorial circular orbit (Abramowicz et al. 1978 ; Bardeen et al. 1972 ; Li Bambi 2013 ). Indeed, for l l ms , the fluid angular momentum is too low and there is no stable accretion disk (the fluid directly plunges onto the black hole), while, for l > l mb , the fluid angular momentum is too high and there is no accretion (there is just a fluid rotating around the black hole). In the presence of an accretion disk, when , the inner edge of the disk, R in , is between the marginally bound orbit and the ISCO, i.e., R mb R in R ISCO (Bardeen et al. 1972 ; Abramowicz et al. 1978 ; Li Bambi 2013 ); that is, for a given black hole spin, the inner edge of a Polish donut disk is always smaller than the inner edge of a Novikov–Thorne disk. Since the exact location of the inner edge of the disk is crucial in reflection spin measurements, we can already suspect that systematic uncertainties related to the use of an incorrect disk model to fit the data of a source can be important. 3. Simulations and Fits In order to estimate the impact of the disk structure on spin measurements with the iron line method, we use the code presented in Riaz et al. ( 2020 ) to calculate the theoretical reflection spectrum of a Polish donut disk. We use the xillver model (García Kallman 2010 ; García et al. 2013 ) to calculate the nonrelativistic reflection spectrum at the emission point in the rest-frame of the gas in the disk, and we use our ray-tracing code to convolve this nonrelativistic spectrum into a relativistic spectrum of a thick disk. At this point, we need to choose the properties of our source and of our instruments to simulate the observation of a black hole with a thick accretion disk. As X-ray mission, we consider NICER (Gendreau et al. 2016 ), which has a good energy resolution near the iron line and is thus suitable for X-ray reflection spectroscopy measurements. We assume that the source has a very simple spectrum, just a power-law component and a reflection component. In XSPEC language (Arnaud 1996 ):where TBABS takes the Galactic absorption into account (Wilms et al. 2000 ), powerlw is the power-law component generated by inverse Compton scattering of thermal photons from the disk off
Paper: 2018ApJ...854L..25P (fulltext)
broadly consistent with the predictions of Fairhurst ( 2011a ). While the inclusion of LIGO-India was not examined in Fairhurst ( 2011a ), we recover a median localization of 9.3 deg 2 for the HIKLV network. 4.3. Comparison with Previous Studies We can compare our result to several studies of advanced network sky localization that have already been performed. We compared our results to those from analytic expectations of Wen Chen ( 2010 ) and Fairhurst ( 2011a ) throughout Sections 4.1 and 4.2 , respectively. The expression used to calculate the expected error regions from Wen Chen ( 2010 ) produces values that are consistent with both Fairhurst ( 2011a ) and our studies. In addition, our study remains broadly consistent with the analytic results of Vitale Zanolin ( 2011 ), recovering similar scalings in localization improvement for successive detector networks for a canonical BNS to our results in Table 2 . Studies such as Singer et al. ( 2014 ) have examined the sky localization capabilities of two- and three-interferometer networks of LIGO and LIGO–Virgo, using more realistic assumptions. In particular, they carried out an end-to-end study of the localizations from the outputs of one compact binary search pipeline used in the first two observing runs (Messick et al. 2017 ). They use a stream of data with colored, but Gaussian, noise and apply a minimum network S/N cut of 12 for each event. Singer et al. ( 2014 ) projected that Virgo would join the LIGO instruments during the second observation run at the sensitivity indicated in Abbott et al. ( 2016b ). They examine a population of BNSs in the HLV configuration, obtaining a median sky localization of 235 deg 2 at the 90% credible level. For our study of NSBH localization, we obtain a median sky localization of 193 deg 2 , seemingly contrary to the expectation that BNS sources are better localized than NSBHs. However, the fiducial networks in our study are have better sensitivities over a given frequency range compared to Singer et al. ( 2014 ), leading to improvements in the effective bandwidth measured by a given instrument. So, our improved localization from that predicted by Singer et al. ( 2014 ) can be attributed to our use of 2018-era instruments as opposed to the less-sensitive 2016 configuration used in Singer et al. ( 2014 ). Rodriguez et al. ( 2014 ) resolve sky localization for a BNS population in which each event realizes a network S/N of 20. On the contrary, our NSBH study samples a wide distribution of network S/Ns, meeting a minimum S/N threshold. 6 6 We require an S/N of at least 5.5 in the second-loudest detector. Therefore, our study probes the S/N regime at which detections are most likely to occur, while Rodriguez et al. ( 2014 ) ignores this low network S/N range. Nevertheless, our Figure 1 is qualitatively comparable to Figures 6 and 7 in Rodriguez et al. ( 2014 ). In particular, the increased isotropy of a three- to four-instrument network and the subsequent,
Paper: 2017ApJ...841..118T (fulltext)
, H β and H α , and narrow [O iii ] λλ 4959, 5007, and Balmer lines. The absorption features around 5450 Å are telluric. The arc spectrum shows some narrow emission lines, such as [O ii ] λ 3728, H β , [O iii ] λ 5007, H α , and [N ii ] λ 6583, and some weak absorption lines, such as Ca ii K λ 3933 and Ca ii H λ 3968. The arc spectrum is contaminated by the QSO, showing broad H α and Mg ii lines (see Section 2.2.2 for details). The redshift of the QSO is measured to be 0.3954 ± 0.0001 from the [O iii ] λλ 4959, 5007 emission lines. The redshift of the arc, using the QSO-contamination-subtracted spectrum, is measured to be 0.3970 ± 0.0001 via [O ii ] and H β emission lines, and consistent with the value of 0.3968 ± 0.0003 measured via the Ca ii K λ 3933 and Ca ii H λ 3968 absorption lines. The arc is likely located nearer than the QSO and moving toward it with a line-of-sight velocity of roughly 340 km s −1 . Figure 4. VLT FORS2 spectrum of the QSO PHL 6625 with model components. We fitted the spectrum from 2750 to 7100 Å, where the S/N is high. The windows used to fit the continuum are marked in black bars, and pixels with telluric absorptions around 5450 Å were excluded in the fitting. The dereddened spectrum is shifted to the rest frame of the QSO at z = 0.3954. Purple: power law; yellow: pseudo-continuum = power law + Fe ii templates; red: summed line model; blue: narrow-line components; green: broad-line components. Figure 5. VLT FORS2 dereddened spectrum (black) of the arc galaxy with the QSO contamination (blue), the galaxy population synthesis model (green), and the total model (red). All emission lines and telluric absorptions are masked during the population synthesis fit. The dereddened spectrum is shifted to the rest frame of the arc at z = 0.3970. 2.2.1. Decomposition of the QSO Spectrum The QSO spectrum is decomposed into multiple emission components following Shen et al. ( 2008 ). The fitting is performed in the rest-frame wavelength range of 2750–7100 Å, where the S/N is sufficiently high. First, a pseudo-continuum was fitted to the spectrum in some continuum windows 12 12 2925–3500 Å, 4200–4230 Å, 4435–4700 Å, 5100–5535 Å, 6000–6250 Å, and 6800–6900 Å at the rest frame. , consisting of a power-law component and Fe ii templates in both the Mg ii region (Salviander et al. 2007 ) and the H β region (Boroson Green 1992 ). The continuum-subtracted line spectrum was then fitted with multiple Gaussian components: three Gaussians for each of the H α , H β , and Mg ii broad components; five Gaussians for narrow lines in the H α region; 13 13 Two for [S ii ], two for [N ii ], and one for narrow H α . three near H β; 14 14 Two for [O iii ] and one for narrow H β . and one for narrow Mg ii . All of the narrow lines are imposed to have the same shift and width. The monochromatic continuum luminosity λL λ of the power-law component at the rest frame 3000 and 5100 Å, the power-law spectral index, and the broad-line width derived from the
Paper: 2020MNRAS.493.3159G (fulltext)
acceleration. Our inference above is qualitatively different from these models in the sense that it is independent of SNe event in a star cluster. We also note that LOFAR observations show PeV CRs to be enriched in low-Z nuclei (Buitink et al. 2016 ), consistent with their being accelerated from CNO enriched wind material in SBs (Thoudam et al. 2016 ). It is also possible that the hard component of CRs due to WTS is the one inferred to be present in molecular clouds (de Boer et al. 2017 ). Therefore, WTSs are potential PeVatrons and they should be studied in detail in this regard. 5.5 Decoupling of grammage from ISM Our proposed scenario can also provide the astrophysical framework for the phenomenologically motivated models of CR propagation in which the grammage traversed by CRs are mostly near the source (Blasi Serpico 2009 ; Cowsik Madziwa-Nussinov 2016 ; Eichler 2017 ; Biermann et al. 2018 ). In these models, the diffusion property of CR particles is assumed to be different inside the ‘cocoons’ surrounding the CR acceleration sites than elsewhere in the ISM in order to alleviate problems associated with secondary production. The scenario of CR acceleration in SBs satisfies the basic premises of these models, since most of the CR collisions occur in the outer shell of shocked ISM, whereas CRs are advected by the wind in the inner regions. The grammage suffered by CRs in the shocked ISM region is ∼10 g cm −2 , comparable to the total grammage of GCR (Gabici et al. 2019 ), since the typical residence time is t r ∼ 1 Myr (∼κ/ v 2 , for a diffusion coefficient κ ∼ 10 27 cm 2 s −1 and v ∼ 50 km s −1 , the typical outer shock speed), and typical density in this region is ∼10 cm −3 . Therefore, CR grammage can be decoupled from Galactic residence time, alleviating a number of outstanding problems (Butt 2009 ; Gabici et al. 2019 ). Instead of listing them all, we mention one of the problems here, that of the observed scaling of light elements with metallicity in halo stars (Parizot Drury 1999 ). In the case of isolated SNRs (which accelerates ISM particles), the Li/Be/B abundance is expected to scale as Z 2 , since both CR and the target gas share the same metallicity. In the case of WTS, the metallicity of CRs is independent of that of the ambient ISM, the abundance of spallation products should scale as Z , as observed (Parizot Drury 1999 ). 5.6 Caveats and future directions In this study, we show the first calculations for estimating the energetics of the wind and SN shocks in isolated stars/star clusters. We have used a shock detection algorithm and three different tracers to quantify the fraction of shock energy processed by wind, SN ejecta, and ambient material. Although in the above sections, we have shown that this study has succeeded in revealing answers to various problems in GCRs, below we highlight some of the limitations of our models and future directions. Non-radiative strong shocks: We have considered the Rankine–Hugoniot
Paper: 2020ApJ...905..135B (fulltext)
effect saturates in the DML, so MOND can provide only a limited amount of additional stability even for a galaxy with arbitrarily low surface density. This might explain why low surface brightness disk galaxies often have spiral features (McGaugh et al. 1995 ), usually taken as a sign that disk self-gravity is important (Lin Shu 1964 ; Das et al. 2020 ). The global stability of disks thus offers a promising way to break the CDM-modified gravity degeneracy in galaxies. M33 is one of the nearest galaxies with a rather low surface brightness. S19 recently conducted isolated N -body and hydrodynamical (hydro) simulations of M33 in the CDM context. As with many other galaxies, its observed RC required them to use a cored DM profile. They found it rather difficult to explain the weak observed bar together with the two-armed spiral (Corbelli Walterbos 2007 ) in the presence of a live CDM halo. Weak bars are also observed in many other galaxies (Cuomo et al. 2019 ). They are problematic in the CDM picture because a bar in the baryons creates a response in the live DM halo, which would cause a resonant effect (Athanassoula 2002 ). This was demonstrated explicitly by S19 , who compared their live halo simulations to a “frozen halo” simulation in which the DM particles provide a fixed gravitational potential. 7 7 This can be thought of as the DM particles providing gravity but not moving. S19 showed that M33 would be more stable in a frozen halo, leading to a weak bar similar to that in the observed M33. Such a model is of course unphysical in the ΛCDM context, but useful to better understand the dynamics. Interestingly, a frozen halo is somewhat similar to gravity theories like MOND where galaxies lack hypothetical DM. This is because features in the disk very rapidly become irrelevant to the potential as one moves away from its plane. Consequently, perturbations in the disk have very little effect on the Milgromian “phantom DM” (PDM) halo because its density depends only on the local potential. 8 8 PDM is the DM density distribution that would be needed for the Newtonian gravitational field to equal the MOND one of the baryons alone. The PDM distribution of thin exponential disks was visualized in, e.g., Lüghausen et al. ( 2013 , 2015 ). Indeed, this is why a bare disk can be stable in MOND. The lack of an actual DM halo automatically removes the possibility of bar–halo angular momentum exchange. However, disk self-gravity is enhanced in MOND compared to the case of a frozen DM halo. The presence of an EFE would slightly decrease this enhancement, but would also decrease the spheroidal potential support from the PDM. In general, it would also break the axisymmetry and up–down symmetry of the problem. Clearly, numerical simulations are needed to properly understand how disk stability is influenced by a fundamental modification to the gravity law. Several MOND N -body simulations have considered isolated disks, starting with the seminal work of Brada
Paper: 2018MNRAS.475.5144S (fulltext)
little variation of A e with rotation. The peak values are also quite different: 150 per cent for HD 37776, 50 per cent for HD 189775, and only 10 per cent for HD 175362. Stars with highly variable A e likely exhibit greater surface abundance anisotropies than stars in which A e is less variable. In order to obtain a single number with which to characterize the average strength of deviations between the magnetic curves derived for different elements for a given star, we used the idl curvefit routine to generate second-or third-order harmonic best fits, which were then integrated as functions of rotational phase. The significance of this integrated value is then A e /σ e , which is given for each star in the second-last column of Table 2 , where σ e is the mean uncertainty in the individual A e measurements. If A e /σ e 1, it can be concluded that differences in 〈 B z 〉 between different elements are a consequence of noise rather than real variations. It is this significance in which we are interested, as it will help to decide whether accurate measurements are available using metallic lines, or whether H lines should be used instead. Fig. 9 shows the histogram of A e /σ e for all stars. For the full sample (top panel), the median significance is 2.7σ. For the sub-sample of He-weak and He-strong Bp stars, the median of the distribution is at 2.9σ (second panel from the top). Amongst pulsating β Cep and SPB stars (third panel from the top) and the remainder of the sample with no apparent peculiarities (bottom panel), only one star, HD 96446, has A e /σ e ≥ 3: as HD 96446 is also a He-strong star, this likely reflects the star's chemical peculiarities. The medians of these distributions are 1.4σ and 1σ, respectively. While there are only eight non-chemically peculiar stars for which A e could be measured, these results are consistent with an origin of the effect in distortion due to chemical spots, and confirm that when chemical peculiarities are not present, there is no significant difference in 〈 B z 〉 when measured using different chemical elements. Even amongst the Bp stars, although there are several stars for which A e /σ e ≫ 3, differences in 〈 B z 〉 measured from different elements are negligible for many of the stars. Figure 9. Significance histogram for A e for (top–bottom) all stars in the sample for which A e could be measured, Bp stars, pulsating stars, and chemically normal, non-pulsating stars. The dashed line marks 3σ significance. Approximately half the sample shows at least a 3σ spread in 〈 B z 〉 measurements conducted with different elements. All of these are Bp stars: the only pulsator with A e /σ e ≥ 3 is HD 96446, which is also a He-strong star. It should be noted that while the mathematical treatment of A e is based upon random uncertainties, these variations are in fact systematic . Unfortunately, it is not obvious what simple mathematical tools might be utilized in order to capture the systematic variations between