adsabs/WIESP2022-NER huggingface dataset:
Paper: 2017ApJ...849...57D (fulltext)
(Spergel et al. 2003 ). All the errors and upper limits are quoted at a 90% confidence level, unless otherwise specified. 2. Data and Catalogs The NuSTAR extragalactic survey program consists of tiered observations of well-known survey fields: (i) a set of deep tiled pointings covering the E-CDFS (Mullaney et al. 2015 ) and EGS (A. Del Moro et al. 2017, in preparation) fields, with an area of and arcmin 2 , respectively, and a total exposure of 1.49 Ms in each of the two fields, reaching a maximum exposure of ≈220 ks in E-CDFS and ≈280 ks in EGS (at 3–24 keV in each focal plane module (FPM), vignetting corrected); (ii) a medium-depth set of 121 tiled pointings covering ≈1.7 deg 2 of the COSMOS field (Civano et al. 2015 ) for a total exposure of 3.12 Ms, with a maximum depth of ≈90 ks in the central ∼1.2 deg 2 ; (iii) a serendipitous survey consisting of all serendipitous sources detected in any NuSTAR targeted observation (excluding all sources associated with the target; Alexander et al. 2013 ; Lansbury et al. 2017 ). This last tier of the survey spans a wide range of depths (see, e.g., Aird et al. 2015a ; Lansbury et al. 2017 ) and has the largest sky coverage, reaching deg 2 to date (still ongoing). Since in this work we require low-energy ( keV) Chandra data, together with the NuSTAR data, to produce broad X-ray band composite spectra, we limit our analyses to the sources detected in the E-CDFS, EGS, and COSMOS fields, which have good Chandra coverage and redshift completeness (spectroscopic and photometric), while we exclude the serendipitous survey sources owing to the heterogeneity of the available ancillary data. Although XMM-Newton observations are also available for the E-CDFS and COSMOS fields, we do not include these data in our analyses because combining the data from different X-ray instruments to produce the composite spectra can cause significant distortions in the resulting spectra (see Section 3.2.1 ). Our sample consists of 182 AGNs (see Table 1 ), 49 detected in E-CDFS (Mullaney et al. 2015 ), 42 detected in EGS (A. Del Moro et al. 2017, in preparation), and 91 detected in COSMOS (Civano et al. 2015 ). Of these sources, 64 (∼35%) are detected in the NuSTAR HB (8–24 keV). 25 25 Most of the sources that are formally HB undetected (∼65%) are actually detected also above 8 keV, but with lower significance than the false-probability threshold adopted by Civano et al. ( 2015 ), Mullaney et al. ( 2015 ), and A. Del Moro et al. (2017, in preparation). Table 1 NuSTAR Source Sample Summary Notes. a Number of HB (8–24 keV) detected sources. b Including both spectroscopic and photometric redshifts; the number of spectroscopic redshifts is reported in parentheses (also in the fifth and sixth columns). In Civano et al. ( 2015 ) and Mullaney et al. ( 2015 ), the NuSTAR sources have been matched to the Chandra and/or XMM-Newton point-source catalogs available in these fields (Lehmer et al. 2005 ; Brusa et al. 2010 ; Xue et al. 2011 ;
Paper: 2019ApJ...884...20B (fulltext)
Nonlinear regression is the best model, according to SIC. The break period varies between 2.35 log P 2.55 days at different wavelengths for all statistical methods. The F-test suggests that the variances of the quadratic and two-slope models are statistically similar. Further, the variances of a quadratic model and a linear relation for P 400 days are also similar in most cases. Therefore, we adopt a break period of 300 days () as the optimal period at which the change in the slope of PL and PLC relations occurs for O-rich Miras. The break period possibly shifts to longer periods (∼350 days, days) at NIR wavelengths. NIR PLRs are discussed in detail in Section 5 . Table 4 Results of the Statistical Tests to Find a Kink in the O-rich Mira PL and PLC Relations in the LMC Note. SIC: Schwarz Information Criterion used to estimate the break period in a two-slope model. RW: Random walk, TM: The Testimator, FT: F-test p ( f ): Probability of acceptance of the null hypothesis that the variance of a quadratic and a linear relation for P 400 days are equal. 4. Stability of Maximum Light for Miras 4.1. Photometric Variability at Max-light Optical photometry provides clear evidence that Miras display significantly smaller scatter in the PLRs at max-light than at mean-light. The pulsation properties and brightness variations of Miras are not strictly periodic and vary from cycle-to-cycle. Therefore, it is important to investigate the stability of max-light over different pulsation cycles. However, apart from OGLE I -band time-series observations, there is a lack of continuous photometry for Miras in the LMC, and even OGLE data do not necessarily cover the epochs of max-light during all pulsation cycles. High-precision photometry and subhour cadence from Kepler can be extremely useful to test the stability of max-light over different pulsation cycles. Bányai et al. ( 2013 ) analyzed variability in M giants from Kepler and identified a dozen long-period Mira-like variables. We estimate the scatter in the max-light magnitudes and median magnitudes for Miras over different pulsation cycles. For OGLE Miras, the sample is restricted to those stars for which GPR predicts max-light magnitudes with a precision better than three times their median uncertainty, and the time coverage is equivalent to the period of each Mira. The top panel of Figure 9 displays the comparison of the scatter in the median- and max-light magnitudes for OGLE and Kepler Mira candidates. When the scatter in both sets of magnitudes is smaller than 0.4 mag, max-light magnitudes seem to be more stable for most O-rich Miras. C-rich Miras often exhibit long-term trends and thus display large scatter both at median- and max-light. The bottom panel of Figure 9 presents the variation in min-light magnitudes versus the scatter in median magnitudes. For a majority of the stars, min-light magnitudes display significantly larger scatter than median magnitudes over multiple pulsation cycles. The
Paper: 2018ApJS..234...14M (fulltext)
analysis the Si q + ions are electrostatically deflected by 45°, then focused and collimated by a three-element einzel lens and a series of three apertures prior to entering the target gas cell. The ions undergo CE collisions with the target gas, lowering their charge states by (considered here) one and two exchanges. Currents for each CE step are measured in a retarding-potential analyzer placed at the end of the CE cell, followed by a deep Faraday beam trap. The target gas-cell pressure is monitored using a calibrated capacitance manometer (CM). Individual charge-state currents are measured through the application of a stepped retarding potential , the value of which depends on the incident charge-state q and ion energy through the relation = , where n is the number of exchanges. Thus, a series of retardations , , , will block the charge states q , q − 1, and q − 2, for a total of two exchanges, the maximum encountered in the present study. As in Machacek et al. ( 2014 ), the cross section σ q , q − n for the exchange of n electrons is given by Here, T is the target gas temperature, k is the Boltzmann constant, P is the gas pressure, L is the effective gas-cell length, and R q,q − n is the particles’ current ratio, given bywhere the currents , , correspond to measured current steps after application of the retarding potentials , , . Additional details on the 28 Si q + mass/charge spectrometry, retarding-potential stepping, measurement of background levels, presence of metastable ion states, tests for gas-cell pressure effects on CE cross sections, and data acquisition, together with sources and the levels of random and systematic errors, are given in Machacek et al. ( 2014 ). The He and H 2 used herein were ultrahigh purity-grade (99.999%) gases, and the 28 Si q + ions were produced by dissociative ionization of silane gas (SiH 4 , UHP-grade 99.999%) in the ECR plasma. 3. Results and Discussion A summary of the sources and magnitudes of the experimental errors is given in Table 1 . Shown in Figures 1 and 2 and listed in Table 2 are results of the present measurements of single and double total cross sections for 28 Si q + ions ( q = 7–10) colliding with He and H 2 at the indicated energies. Also given are the statistical error limits at the 2 σ level of confidence. In Table 2 note that the ratio of single-to-double CE cross sections for 28 Si q + + He is approximately 4.1 ± 1.2 (1 σ ) for all the charge states, with the exception of 28 Si 10+ at 2.25 keV/u, where the double-exchange cross section is small (0.104 × 10 −17 cm 2 ). The single-to-double ratio for 28 Si q + + H 2 shows larger variations, from 0.04 to 49.5. The reason for this is not understood. In order to provide some theoretical insight, calculations of the single-exchange CE cross sections were carried out in the static one-electron model (Mann et al. 1981 ). Here, the captured electrons mainly populate states with principal quantum number n that depend on the charge q of the
Paper: 2018ApJS..236...31K (fulltext)
velocity, which was mostly determined by the NH 3 line observations (Table 1). Figure 5.50. 3775 The detected 22 GHz H 2 O maser spectra of the sources detected only at 22 GHz. Red and Black colors represent the first and second epochs, respectively. The source name is given in the top-left corner of the upper panel and “M” is attached for the source associated with 6.7 GHz CH 3 OH maser emission (see Section 3.1). The observing dates are shown in the top-right corner. The vertical dotted line indicates the systemic velocity, which was mostly determined by the NH 3 line observations (Table 1). Figure 5.51. 3816 The detected 22 GHz H 2 O maser spectra of the sources detected only at 22 GHz. Red and Black colors represent the first and second epochs, respectively. The source name is given in the top-left corner of the upper panel and “M” is attached for the source associated with 6.7 GHz CH 3 OH maser emission (see Section 3.1). The observing dates are shown in the top-right corner. The vertical dotted line indicates the systemic velocity, which was mostly determined by the NH 3 line observations (Table 1). Figure 5.52. 3866 The detected 22 GHz H 2 O maser spectra of the sources detected only at 22 GHz. Red and Black colors represent the first and second epochs, respectively. The source name is given in the top-left corner of the upper panel and “M” is attached for the source associated with 6.7 GHz CH 3 OH maser emission (see Section 3.1). The observing dates are shown in the top-right corner. The vertical dotted line indicates the systemic velocity, which was mostly determined by the NH 3 line observations (Table 1). Figure 5.53. 3878 The detected 22 GHz H 2 O maser spectra of the sources detected only at 22 GHz. Red and Black colors represent the first and second epochs, respectively. The source name is given in the top-left corner of the upper panel and “M” is attached for the source associated with 6.7 GHz CH 3 OH maser emission (see Section 3.1). The observing dates are shown in the top-right corner. The vertical dotted line indicates the systemic velocity, which was mostly determined by the NH 3 line observations (Table 1). Figure 5.54. 3880 The detected 22 GHz H 2 O maser spectra of the sources detected only at 22 GHz. Red and Black colors represent the first and second epochs, respectively. The source name is given in the top-left corner of the upper panel and “M” is attached for the source associated with 6.7 GHz CH 3 OH maser emission (see Section 3.1). The observing dates are shown in the top-right corner. The vertical dotted line indicates the systemic velocity, which was mostly determined by the NH 3 line observations (Table 1). Figure 5.55. 3914 The detected 22 GHz H 2 O maser spectra of the sources detected only at 22 GHz. Red and Black colors represent the first and second epochs, respectively. The source name is given in the top-left corner of the upper panel and “M” is attached for the source associated with 6.7 GHz CH 3 OH
Paper: 2019MNRAS.485.5210K (fulltext)
can form a favourable overlap elsewhere on the backbone than at C2. Some electronic redistribution is required so it may not be as favourable for H to attach to C3 or N as it is at C2. This can account for the lower stability of the ketenimine product and the observed ∼9:1 ratio of the two products shown in Table 1 . We explored these two pathways in the related but simpler system H 2 CCN + H. Replacing the methyl group of CH 3 CHCN with H should not greatly affect the energetics of H-addition. Like CH 3 CHCN, the H 2 CCN radical shows resonance behaviour between two forms, as indicated by radical character on all three heavy atoms. For H 2 C1-C2-N, the Mulliken spin populations at the UMP2/aug-cc-pVTZ level on C1, C2, and N are 1.09, −0.64. and 0.65, respectively, which are quite similar to the spin densities on the analogous atoms in CH 3 -CH-CN above (there are also minor amounts of spin density on the two H atoms). We might therefore expect to form both CH 3 CN (methyl cyanide) and H 2 CCNH (ketenimine) by adding H to C1 and N of H 2 CCN, respectively. To test the likelihood of forming both products, we ran additional calculations. These calculations required a multireference treatment since closed-shell singlet molecules are being formed from open-shell doublet reactants. At the MRCI + Q/aug-cc-pVTZ level, we found that neither addition involves a reaction barrier. The issue thus becomes trying to understand why there is a preference for the cyanide product over the ketenimine product. With no barriers for either addition, the favourability of one product over the other must be due to the specific character of the full reaction surface. In addition of H to CH 3 CHCN, ethyl cyanide is more stable than the ketenimine product, and the same is true when adding H to H 2 CCN: methyl cyanide is more stable than ketenimine. The first channel has a ‘wider’ entrance on the potential energy surface; it is deeper and extends further out from C1. This is likely what favours cyanide formation over ketenimine formation in these reaction systems. To gain further insight into the ratio of the two products, we note that the study of delocalized radicals containing nitrogen atoms has been theoretically investigated in the case of the simplest of the cyanoalkyl radicals, CH 2 CN. Delbecq ( 1983 ) has shown using valence-bond analysis that cyanomethyl radical CH 2 CN has two different limit structures, C°H 2 CN and CH 2 = C = N° with relative weights of 83 per cent and 14 per cent, respectively. The H 2 CN π-delocalized structure was found with a low weight of the order of 3 per cent. Similarly, under our experimental conditions, the formation of CH 3 − C°H − C≡N through CH 2 = CH − C≡N + H reaction, would lead as shown in Fig. 8 , to the formation of a second localized form with ketenimine character CH 3 − CH = C = N° that explains the possibility of reaction at the nitrogen end to form CH 3 CH = C = NH. Figure 8. Two different limit structures and one
Paper: 2017ApJ...845...32K (fulltext)
of the relatively isolated star formation. External compression of the core (e.g., Frau et al. 2015 for the possibility of the cloud–cloud collision in the Pipe Nebula) may also initiate the star formation in the core. Rotation of the core does not affect its stability. The rotational kinetic energy of the core is only approximately 1% of the gravitational energy (Aguti et al. 2007 ). Uncertainty arising from the use of the plane-of-sky field component increases the importance of magnetic support. The magnetic field close to a pole-on geometry leads to the condition (subcritical). Therefore, 3D modeling of the polarization structure is desirable in order to determine the line-of-sight inclination angle of the core’s magnetic axis. The formation of a nearly critical core accompanied by a distorted magnetic field is an open problem. Consideration of the interplay among gravity, thermal pressure, and turbulence is necessary in order to understand the formation and evolution mechanism of FeSt 1–457. Quasi-static contraction under the support of a magnetic field (Shu 1977 ) requires a strong magnetic field, and this is not the case for FeSt 1–457 (, magnetically supercritical). In the other extreme, supersonic turbulence can produce cores that collapse dynamically and that are accompanied by highly supersonic infalling motion (e.g., Mac Low Klessen 2004 ). This is also not the case for FeSt 1–457 because nearly critical stability was shown in the present study and radio observations (Kandori et al. 2005 ; Aguti et al. 2007 ) showed quiescent kinematic gas motion in the core. Core formation between the two extreme models may account for the formation of FeSt 1–457 (e.g., Nakamura Li 2005 ; Basu et al. 2009a , 2009b ). Further theoretical studies on such a case are desirable. The relative importance of the magnetic field for the support of the core is investigated through the ratios of the thermal and turbulent energy to the magnetic energy, which are expressed as and , where , , and denote the isothermal sound speed at 9.5 K, the turbulent velocity dispersion, and the Alfvén velocity, respectively. The numerical coefficients differ in the equation depending on whether the velocity dispersion is in 1D or 3D. These ratios were found to be and . FeSt 1–457 is dominated by a static magnetic field and thermal support with a small contribution from turbulence. Turbulence could be an important parameter in the formation and development of the core, although it seems to have already mostly dissipated at the observed stage in the prestellar evolution. Therefore, a static magnetic field, the predominant component of anisotropic pressure in the core, should act as a key parameter controlling the star formation process in FeSt 1–457. Figure 8 shows the relationships among the orientations (position angles) of the axis of the elongation of the core (), the rotation axis (, Aguti et al. 2007 ), and the magnetic field axis (). Note that the elongation of the
Paper: 2018MNRAS.476.5216D (fulltext)
of SEGUE dwarfs. They measured, for the disc as a whole, a vertical metallicity gradient of dex kpc −1 , which is in agreement with our measurement. However, their low-α population has a gradient consistent with zero: d[M/H]/d dex kpc −1 . However, their intermediate α sub-population with 0.2 [α/M] 0.3 has d[M/H]/d dex kpc −1 , which agrees with our low-α metallicity gradient. We thus conclude that the discrepancy between our result and that of Schlesinger et al. ( 2014 ) is largely due to the chemical separation criteria (also see Ciucă et al. 2018 , who found an age-dependence for the thin disc vertically metallicity gradient, such that the youngest population has a flatter gradient). Mikolaitis et al. ( 2014 ) measured a slightly shallower gradient of d[M/H]/d z = −0.11 ± 0.01 for Gaia -ESO dwarfs and giants. The Gaia -ESO sample is more metal-poor overall, and Mikolaitis et al. ( 2014 ) separated thick disc stars by the location of underdensities in their [Mg i /M] histograms (their fig. 3). The dividing line is at different values of [Mg i /M] for different metallicity regimes. In fig. 10 of Mikolaitis et al. ( 2014 ), it is clear that their sample is biased against metal-rich stars, such that there are very few stars with [M/H] > 0 (also see Stonkutė et al. 2016 ). 5.2.2 The high-α population The vertical metallicity distribution of the thick-disc (high-α) stars is relatively flat compared to the low-α population, at −0.058 ± 0.003 dex kpc −1 . Several authors have measured the vertical gradient for the thick disc, using different methods to define this population. Earlier studies, such as Gilmore et al. ( 1995 ) and Prieto et al. ( 2006 ), reported no vertical metallicity gradient in the thick disc (Prieto et al. 2006 quoted an upper limit of d[M/H]/d z = 0.03 dex kpc −1 ). More recently, Boeche et al. ( 2014 ) concluded that the vertical metallicity gradient of the thick disc is consistent with zero, based on a sample of RAVE giants. However, other studies, using a combination of metallicity or kinematics to separate the thick disc, have reported a shallow metallicity gradient. Katz et al. ( 2011 ) observed sub-giants at two lines of sight: ( l, b ) = (51°, 80°) and (5°, 46°) at low resolution. Their MDFs show signs of bimodality, and the thick disc was defined as stars centred around [M/H] ≈ −0.5 dex. The vertical gradient measured by Katz et al. ( 2011 ) is −0.068 ± 0.009 dex kpc −1 , consistent with our value. Ruchti et al. ( 2011 ) observed a number of metal-poor thick disc candidates at high resolution using the MIKE, FEROS and UCLES spectrographs (λ/Δλ ≈ 35 000–45 000). They classified their stars based on a Monte Carlo simulation of space motion U, V, W , assuming Gaussian errors on the velocities and distances. By further restricting their α-enhanced sample with thick disc kinematics to metal-poor stars only ([M/H] ≤ −1.2), they avoid most thin disc contamination. The measured gradient is −0.09 ± 0.05 dex kpc −1 ,
Paper: 2021ApJ...914...16H (fulltext)
DR2 is designed to observe stars in the Milky Way, it does not always provide accurate measurements of sources that are extended or in very crowded regions. We urge caution and follow-up observations of the objects in Table 5 , especially those within 10′ of the center of NGC 5128. B.2. Conflicting Measurements Thirty-two of the identified GCs in NGC 5128 have conflicting classifications in the literature. A majority of these conflicting classifications occur because the target was measured to have a radial velocity within the standard acceptable range of ≈250–1000 km s −1 by one group but outside this range by another. Additionally, two objects were identified as GCs based on their radial velocities but later found to be stars based on HST/ACS imaging (see note 9 in Table 6 ). Another object was subsequently identified as a star because of its spectrum (see note 5 in Table 6 ). These objects with conflicting measurements have been matched between studies based on position and, in many cases, by their ID as well. When examining these objects with PISCeS data, several of them can be visually identified as background galaxies. We believe seven of these objects are indeed true GCs because a majority of their measurements classify them as such and they do not visually look like galaxies; thus, they are included in Table 4 . For each object with conflicting measurements, Table 6 first lists the R.A. and decl. in degrees (J2000). The position measurements are from the PISCeS data set to provide a uniform astrometric system. Next, we list the conflicting radial velocities and errors (kilometers per second) along with their sources, if applicable. We also include our classification of each object, as well as further information indicated by the “Notes” column. Many sources experience a discrepancy between Woodley et al. ( 2007 ) and Beasley et al. ( 2008 ). It appears that between sharing data with the Woodley group and publishing their own paper, the Beasley group updated measurements on several targets. We note that one of these objects has a mismatch in both its velocity measurement and coordinates. Woodley et al. ( 2007 ) gave an incorrect position for HH-048 as α = 20068900, δ = −43119111, with a radial velocity of 856 ± 56 km s −1 , citing Beasley et al. ( 2008 ) as the reference for this radial velocity measurement. Beasley et al. ( 2008 ), however, listed HH-048 as a background galaxy ( v r = 16,000 km s −1 ) with the same R.A. but a decl. of δ = −43011917 (note the single difference in the digit following the decimal point). Based upon visual inspection, the position listed in the Beasley catalog is correct and listed in Table 6 . Appendix C An Object with Too Many Names During our literature review, we found an instance where two confirmed GCs with different names shared the same position. The objects identified most recently as GC 0462/T17-1560 and GC 0200/T17-1561 have the same position in the Sinnott et al. ( 2010 ) catalog and positions