adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016A&A...586A.140P (fulltext)
that gas outside halos, not necessarily following any given profile, also contributes to the signal. Furthermore, we would like to conduct an analysis that is as model independent as possible, and thus blind to any specific model for the gas spatial distribution. The correction by a redshift-averaged quantity introduced in the last equation above was justified by Hand et al. (2012) through the need to correct for possible redshift evolution of the tSZ signal in the sources. In practice, this correction allows us to measure relative changes in temperature anisotropies between nearby galaxy pairs after minimizing other noise sources, such as CMB residuals. Fig. 1 Absolute value of versus distance to the observer after obtaining AP kSZ temperature estimates on the real positions of CGs (black line) and on a rotated CG configuration (red line). The green line provides the theoretical expectation for the rms of , and the blue histogram displays the radial distribution of the CGs (see text for further details). We next conduct an analysis into the motivation behind this -dependent monopole correction. In Fig. 1 we display the amplitude of for versus the comoving distance to the observer . The value of the aperture chosen in this exercise is 8 arcmin. The black line refers to the real position of the CGs, while the red line refers to a rotated-on-the-sky configuration of the CGs. Since we are plotting absolute values of , the solid (dashed) parts of these lines refer to positive (negative) values of . The amplitude of these curves is close to the solid green line, which corresponds to the theoretical prediction for the rms of under the assumption that the AP temperature estimates are dominated by CMB residuals. That is, the amplitude of the green curve equals , where K is the rms of the 8 arcmin AP kSZ estimates and is the number of CGs effectively falling under the redshift Gaussian window given in Eq. ( 5 ). While K is computed from real data, its amplitude is very close to the theoretical predictions of Fig. 6 in Hernández-Monteagudo et al. (2006b) , for an angular aperture of 8 arcmin after considering the CMB exclusively. We can see that the amplitude of is lowest for those distances where the number density of CGs is highest (as displayed, in arbitrary units, by the blue histogram in Fig. 1 ). Therefore Fig. 1 shows not only that AP kSZ estimates are dominated by CMB residuals, but also that the difference in for CG pair members lying 50–100 Mpc away in radial distance typically amounts to few times K. We shall show below that this amplitude is not completely negligible when compared to the typical kSZ pairwise momentum amplitude between CGs at large distances, and thus subtracting becomes necessary. However, this should not be the case when cross-correlating AP kSZ measurements with estimates of radial peculiar velocities (see Sect. 3.2 below), since in this case the residual will not be correlated with those velocities and should not
Paper: 2017ApJ...842..123F (fulltext)
that all star formation within Upper Sco has ceased. According to de Geus ( 1992 ) and Preibisch Zinnecker ( 1999 ), star formation in Upper Sco may have been triggered by a supernova-driven shock wave from a massive star in the parent Sco-Cen OB Association (see also Preibisch Zinnecker 2007 ; Slesnick et al. 2008 ). In this context, smooth star formation could occur over a long interval, followed by a supernova that triggers a burst, and subsequent quenching of star formation. 4.1. Comparing Synthetic Populations to the Known Upper Sco Members Our simulations of the star formation history of a cluster correspond to smooth star formation with no trigger (the Smooth-Quench model) and smooth star formation followed by a supernova trigger (Smooth-Burst model). Figure 11 compares the results of these simulations to the stellar population in Upper Sco for the Smooth-Burst model with parameters , , and A = 60, and for the Smooth-Quench model with and . Figure 11. The inferred age vs. temperature of stars in Upper Sco (black, standard deviation in gray dotted lines), compared to results from the Smooth-Quench (top) model, with and , and the Smooth-Burst model (bottom) with A = 60, , and . The shaded regions are the standard deviation from 1000 simulations. The dashed line shows the same results, but when all binaries are resolved. The Smooth-Quench model exhibits a temperature-dependent age measurement up to ∼6 Myr between KM stars and F stars, thereby accounting for half of the ∼12 Myr discrepancy (see Table 2 ). This simplest model demonstrates that an age spread of more than a few Myr in a young ( Myr) population will lead to a significant temperature dependence in age estimates. This qualitative result is required by the evolution of intermediate-mass stars from the Hayashi track to the Henyey track. Cool stars are young and have a median age that changes slowly. For stars with , the median age increases rapidly. At temperatures higher than , stars tend to be older and have median ages that change slowly. These simulation results are consistent with the age spreads of 4–5 Myr over a relatively large temperature range of cool stars (e.g., Herczeg Hillenbrand 2015 ), and with the ∼3 Myr age spread for G type and F type stars (Pecaut et al. 2012 ). Figure 11 shows that the Smooth-Burst model may explain entirely the empirical temperature-age relationship of Upper Sco. The burst increases the population of cool stars at the end of the star formation process, thereby maximizing the possible differences in ages between cool and hot stars. Some of the age difference between K-M type and F type stars may be attributed to the incomplete physics in the calculations of pre-main sequence evolution of low-mass stars (e.g., Feiden 2016 ), although these explanations may introduce problems into other clusters. However, the age difference between G type stars (9 ± 3 Myr) and F type stars (13 ± 1 Myr), as calculated by Pecaut et al. ( 2012 ), are not easily
Paper: 2019ApJ...882L...1F (fulltext)
the PCs would contain the changes. Materials that are detached from the surface, such as the surface-independent clouds in the case of the Earth, would also appear in one of the PCs if they have a large influence on the light curve. If these materials can be constrained by the reflection spectra, their fractions can be derived from the magnitude of the corresponding PC. Once the surface information is extracted from light curves, the surface map of the exoplanet can be recovered from the observational geometry without making any spectral assumptions. Aside from orbital elements, which can be determined from light curve observations, the only two geometry assumptions required for constructing the 2D surface map are the summer/winter solstice and the obliquity. The rotation period of a clear or partially cloudy exoplanet can be inferred from the power spectrum of PCs using a Fourier transform (Figure 3 (d)), which requires the observation frequency to be higher than that of the exoplanet’s rotation. Studies have been performed to identify a planet’s rotational period from light curves at different viewing geometries (e.g., Pallé et al. 2008 ). The summer/winter solstice would coincide with the maxima and minima of the PC time series as long as the asymmetry between the northern and southern hemispheres is noticeable, and the reflection changes monotonically with latitude when the sub-stellar point is near extremum. For the Earth, the peak of the time series of PC2 takes place on 2016 June 15 (Figure 3 (c)), which is within one week of the true value. The obliquity of an exoplanet could be derived through its influence on the light curves. A number of recent publications (e.g., Kawahara 2016 ; Schwartz et al. 2016 ) developed methods for deriving the obliquity using its influence on the amplitude and frequency of light curves. Although these inversion methods are mostly based on a cloudless Earth and known surface spectral features, our SVD analysis of separating clouds from the surface could fill the gap. Some issues exist in constructing and interpreting the retrieved Earth surface map. Degeneracy resulting from the convolution between pixel geometry and spectrum is the dominant factor affecting map construction quality. As discussed in Cowan Strait ( 2013 ) and Fujii et al. ( 2017 ), light curves only cover a small portion of the PC plane, which results in a trade-off between the spatial and spectral variation. Although only PC2 is used for constructing the surface map in this work, its time series covers only a small range of all valid values, whose corresponding land fractions are between 0 and 1. Therefore, we introduce a regularization parameter, λ , when constructing the map, to constrain the pixel values of PC2 within a reasonable range; the resulting effect is described in Appendix B . We select the value of λ based on the ground truth of Earth’s surface map. Glints, features that are small in area but contribute significantly to the
Paper: 2018MNRAS.475.3029T (fulltext)
arcmin away, in a low density environment (Cernicharo et al. 1998 ; Rosado et al. 1999 ; Yusef-Zadeh, Biretta Wardle 2005b ). Evidence of shocked gas associated with core TC1 has also been provided by Yusef-Zadeh, Biretta Geball ( 2005a ). Finally, Rho et al. ( 2004 ) reported the detection of 304 X-ray sources with the Chandra satellite, 70 per cent of which having near-infrared (near-IR) counterparts. X-ray emission was detected from within condensation TC1 and possibly TC4, two star-forming cores with bipolar wings and associated Class 0 candidates (Rho et al. 2004 ). A recent comprehensive review of the Trifid star-forming region has been presented by Rho et al. ( 2008 ). In order to study in detail the principal star-forming cores within M 20, we present here new far-IR images from the Herschel Infrared Galactic Plane Survey (HI-GAL; Molinari et al. 2010 ), combined with a new set of deep near-IR images obtained through broad- and narrow-band filters from the Calar Alto 3.5-m telescope. In addition, archival Spitzer Infrared Array Camera (IRAC) images (Fazio et al. 2004 ; Werner et al. 2004 ), part of its Galactic Legacy Infrared Midplane Survey Extraordinaire (GLIMPSE) survey (Benjamin et al. 2003 ), are used to supplement these observations. The observations are described in Section 2 , whereas in Section 3, we present the results and discussion. Finally, Section 4 lists our conclusions. 2 OBSERVATIONS AND DATA ANALYSIS 2.1 HI-GAL images and photometry The Trifid nebula has been observed in the far-IR as part of the Herschel satellite (Pilbratt et al. 2010 ) HI-GAL survey at 70,160, 250, 350, and 500 μm. The images were reduced using the HI-GAL standard pipeline (Traficante et al. 2011 ). From these, source extraction and photometry were performed using the CuTEx package (Curvature Threshold Extractor), independently at each band (Molinari et al. 2016 ). Fig. 1 (lower right-hand panel) illustrates the Herschel three-colour image of an area of 12 × 12 arcmin 2 around the Trifid nebula, composed from the 70 μm (blue), 160 μm (green), and 250 μm (red) single images. Figure 1. 12 × 12 arcmin 2 (7.0 × 7.0 pc) colour-coded images of the Trifid nebula. Lower right-hand panel: Herschel at 70 μm (blue), 160 μm (green), and 250 μm (red); lower left-hand panel: Spitzer /IRAC at 3.6, 4.5 and 5.8 μm; upper left-hand panel: CAHA 3.5m in the JHK s bands; upper right-hand panel: a visible public image from the European Southern Observatory. All images have the same scale and are centred at the same position. The white contours are from the Herschel 160 μm frame. The white labels indicate the TCs that were studied in this work. The yellow arrows indicate the location of the early-type trapezium system HD 164492, which ionizes the H ii region, and the Herbig–Haro jet HD 399 associated with TC2. The brightest dense cores have been associated with those found at millimetre wavelengths (Lefloch et al. 2008 ) as labelled in Fig. 1 . In addition, filamentary
Paper: 2019ApJ...876..136J (acknowledgments)
We are grateful to J. Chen, J. Delabrouille, Z. Huang, J. Liu, X. Tong, L. Xue, D.-G. Wang, Z. Wang, and S.-Y. Zhou for valuable discussions. We also thank the anonymous referee for constructive comments. Y.F.C. and J.J. are supported in part by the NSFC (grants No. 11653002, 11722327, 11421303, J1310021), by the National Youth Thousand Talents Program of China, by the CAST Young Elite Scientists Sponsorship Program (2016QNRC001), and by the Fundamental Research Funds for the Central Universities. Q.Y.L. is supported in part by the Department of Physics Astronomy at the University of Pennsylvania. Y.Z. is supported in part by the NSFC (Nos. 11675165, 11633001), the CAS Strategic Priority Research Program “The Emergence of Cosmological Structures” (No. XDB09000000) and SRFDP. All numerical computations are operated on the computer clusters “ Linda Judy ” in the particle cosmology group at USTC.
Paper: 2016MNRAS.457.3096A (fulltext)
to 1.99 × 10 −18 cm 3 s −1 for temperatures ranging from 300–14 000 K, respectively, and can be approximated (within 7 per cent) as for temperatures between 300 and 1000 K and for temperatures higher than 1000 K. Because we do not take into account the fine structure of the phosphorus atom (∼227 cm −1 ), our predicted rate coefficients are uncertain at temperatures lower than about 327 K. Moreover, the B 2 Σ + state has a barrier against dissociation into P( 4 S) + O( 3 P). This has already been seen in the theoretical works of Liu et al. ( 2013 ) and Izzaouihda et al. ( 2014 ). Due to the potential barrier, the rate coefficients decrease rapidly at low temperatures. However, the height of the barrier is about 2416.49 cm −1 at 4.8 a 0 , which is about 508.66 cm −1 lower than the value reported by Liu et al. ( 2013 ). Also, the barrier in our PO B 2 Σ + potential appears to be significantly lower than the one shown in fig. 1 of Izzaouihda et al. ( 2014 ). The uncertainty in the barrier height results in a significant uncertainty in the rate coefficients at low temperatures. Other sources of error at low temperatures are quantum mechanical effects such as tunnelling and shape resonance (Bain Bardsley 1972 ; Smith 1989 ; Gustafsson 2013 ). At 300 K, quantum mechanical calculations (Antipov et al. 2009 ; Franz, Gustafsson Nyman 2011 ) show that the classical calculations (Dalgarno, Du You 1990 ; Singh et al. 1999 ; Singh Andreazza 2000 ) underestimate the rate coefficients by about 25 per cent and one order of magnitude for CN (Singh Andreazza 2000 ) and CO (Dalgarno et al. 1990 ; Singh et al. 1999 ) formation, respectively. At temperatures higher than 1000 K, the quantum mechanical (Antipov et al. 2009 ; Gustafsson 2013 ) and semiclassical (Dalgarno et al. 1990 ; Singh et al. 1999 ; Singh Andreazza 2000 ) rate coefficients are in good agreement. In these studies, CN and CO formation was assumed to occur mainly through the A 2 Π and A 1 Π states, respectively. The former potential curve has no barrier, while the latter shows a potential hump. In semiclassical studies, Dalgarno et al. ( 1990 ) and Singh et al. ( 1999 ) used the same set of potentials for the A 1 Π and X 1 Σ + states of CO and different transition dipole moment functions. The rate coefficients obtained by these authors are in agreemnet with each other. However, in quantum mechanical calculations, Franz et al. ( 2011 ) used a different set of potentials with a low barrier height for the A 1 Π state of CO. For CO, the resonance contribution to the rate coefficients arises from quantum mechanical tunnelling through the barrier (Franz et al. 2011 ). At low temperatures, both quantum mechanical and semiclassical theories depend strongly on barrier height. As previously mentioned, the uncertainty in barrier height results in a significant uncertainty in the rate coefficients at low temperatures. Considering all these effects, mainly the uncertainty in barrier height, it is likely that
Paper: 2020ApJ...896L..24D (fulltext)
of orders of magnitude (Dong et al. 2018b ), thereby potentially attaining values as high as ∼0.1 mJy for weak magnetic fields. 3.2. Atmospheric Escape Next, we direct our attention to the central issue of atmospheric ion escape facilitated by the stellar wind in conjunction with the stellar radiation. Figure 2 depicts the pressure balance arising from the stellar wind–planet interactions. The first and second columns present the unmagnetized Venus-like and magnetized Earth-like cases, respectively. As a result of the supermagnetosonic nature of the stellar wind at TOI-700 d, bow shocks are formed ahead of this planet, regardless of the inclusion of the planetary magnetic field. Compared to the unmagnetized Venus analog, the presence of the global magnetic field results in a larger interaction cross-section with the stellar wind for the magnetized Earth-like case. An inspection of the bottom panels of Figure 2 reveals that the planetary magnetosphere (or induced magnetosphere for unmagnetized cases) is compressed, and the total pressure becomes intensified due to the stronger stellar wind pressure. Figure 2. Logarithmic scale contour plots of the total pressure ( P total , units of nPa)—the sum of the ion ( P ion ) and electron ( P electron ) pressure—in the meridional plane based on the stellar wind conditions at P min (first row) and P max (second row). The first column shows the unmagnetized Venus-like case, while the second column depicts the magnetized Earth-like case. The salient results concerning atmospheric escape are depicted in Figure 3 , which shows the calculated oxygen ion density with the associated magnetic field lines in the meridional plane for all six cases. This figure yields some general conclusions. To begin with, as seen from Table 1 , we note that O + constitutes the dominant ion species undergoing escape for all configurations considered here. The atmospheric oxygen escape rates vary from s −1 to s −1 , indicating that they are higher by a few orders of magnitude than the typical escape rates of ) s −1 for the terrestrial planets in our solar system (e.g., Dong et al. 2017b , 2018c ). Figure 3. Logarithmic scale contour plots of the O + ion density (units of cm −3 ) with magnetic field lines (in white) in the meridional plane based on the stellar wind conditions at P min (first row) and P max (second row). The first column shows the unmagnetized Venus-like cases, whereas the second and third columns depict the unmagnetized and magnetized Earth-like cases. Let us first compare the unmagnetized cases, namely, the first and second columns. We find that the Earth-like case exhibits a stronger and broader flux of escaping O + compared to the Venus-like case; this result is also consistent with the atmospheric ion escape rates shown in Table 1 . The reason chiefly stems from the fact that the upper atmosphere of a Venus analog is cooler than its Earth-like counterpart due to the efficient CO 2 cooling caused by 15 μ m
Paper: 2017MNRAS.470.1453Q (fulltext)
(some works on this topic are Uitenbroek 1989 ; Uitenbroek Bruls 1992 ; Belluzzi Trujillo Bueno 2014 ; del Pino Alemán, Casini Manso Sainz 2016 ). However, in the case of the K i lines it seems that this effect is very weak even at extremes heliocentric angles (see the results for the K i D 1 line presented in Uitenbroek Bruls 1992 ). Still, we study the PRD effects for both K i lines using the FALC model [in the previous work the VALC model (Vernazza, Avrett Loeser 1981 ) was used instead]. In this regard, we plot in Fig. 3 the results for the K i D 1 (left) and D 2 (right) lines. For the latter case, we also included the neighbouring photospheric Fe i line (see label 4 in Table 1 ) as it is relatively close. We can see that, for all the selected heliocentric angles, the differences between the computed PRD and CRD profiles are negligible (see the bottom part of each plot). This is in agreement with the results of Uitenbroek Bruls ( 1992 ) and indicates that PRD effects are very small for these lines, at least for the intensity spectrum. Figure 3. PRD (solid) and CRD (squares) profiles for the K i D 1 D 2 lines (we also included the neighbour Fe i line for the latter case). Colours indicate the selected heliocentric angle (larger μ corresponds to an LOS closer to the disc centre and to higher continuum intensity). In addition, we show the difference between PRD and CRD profiles in the bottom of each panel. 3.2 NLTE effects In previous works related to the alkali lines as in Bruls et al. ( 1992 ) it has been demonstrated that the K i D 1 line, in particular its line core, is sensitive to NLTE effects. In other words, the LTE assumption is unable to reproduce the observed line depth. In this section, following the steps of the previous one, we check if the same occurs for the K i D 2 line and also revisit the studies of the K i D 1 line with the FALC model. In addition, as the results of the previous section revealed that PRD effects are negligible for the K i lines, we also compute the differences between NLTE and LTE for the polarization spectrum. We show in Fig. 4 the comparison between an NLTE and an LTE computation for the K i lines (the Fe i is always computed in LTE) assuming a heliocentric angle μ = 1. If we look at the first column, we can see that the line core in case of the NLTE computation (pink) is much deeper than for the LTE case (black). This happens for both potassium lines; the K i D 2 line is very deep with a line core intensity (in NLTE) of almost 0.1 of I c . Moreover, if we examine the rest of the columns, we also find differences between the NLTE and LTE polarimetric signals, the latter always with less amplitude. This is an interesting result that points out the necessity of an NLTE treatment for both the intensity and the polarization spectrum. Finally, we can see that these differences also affect additional spectral features as the Stokes V asymmetries. We plot in Fig. 5 the Stokes V LTE and NLTE profiles for both K
Paper: 2016A&A...592L...5W (acknowledgments)
P.N. acknowledges the support of the Royal Society through the award of a University Research Fellowship, the European Research Council, through receipt of a Starting Grant (DEGAS-259586) and the support of the Science and Technology Facilities Council (ST/L00075X/1). N.B. acknowledges funding from the European Union Seventh Framework Programme (FP7/2007-2013) under grant agreement No. 312725. L.D. and S.J.M. acknowledge support from the European Research Council Advanced Investigator grant, COSMICISM and Consolidator grant, cosmic dust. The H-ATLAS is a project with Herschel, which is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA. The H-ATLAS web site is http://www.h-atlas.org/ . Funding for SDSS-III has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, and the US Department of Energy Office of Science. The SDSS-III web site is http://www.sdss3.org/ . SDSS-III is managed by the Astrophysical Research Consortium for the Participating Institutions of the SDSS-III Collaboration including the University of Arizona, the Brazilian Participation Group, Brookhaven National Laboratory, Carnegie Mellon University, University of Florida, the French Participation Group, the German Participation Group, Harvard University, the Instituto de Astrofisica de Canarias, the Michigan State/Notre Dame/JINA Participation Group, Johns Hopkins University, Lawrence Berkeley National Laboratory, Max Planck Institute for Astrophysics, Max Planck Institute for Extraterrestrial Physics, New Mexico State University, New York University, Ohio State University, Pennsylvania State University, University of Portsmouth, Princeton University, the Spanish Participation Group, University of Tokyo, University of Utah, Vanderbilt University, University of Virginia, University of Washington, and Yale University.