adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016ApJ...833..235S (fulltext)
in Figure 2 . The black solid curves in panel (a) are the same isochrones as used in Figure 2 , but for a distance of 29 kpc. For the Sgr Stream, solid blue lines are used in panel (a) instead of dashed blue lines to represent the Sgr bright arm, and the distance of 23 kpc was adopted for these isochrones. In panel (b), we also show the Law Majewski ( 2010 ) N -body model particles in right- (leading) or left-facing (trailing) triangles. These particles are color coded following Law Majewski ( 2010 ), such that different colors represent the time at which a given debris particle became unbound from the Sgr dSph: magenta for 1.5–3 Gyr ago; and cyan for 3–5 Gyr ago. The gray cross is the same as in Figure 2 . An expectation value of 3 MW halo stars is predicted within the area spanned by the gray cross. The ORPH-F2 field is located on the main Sgr Stream (or the bright arm) at Λ ⊙ = 2264, and we expect Sgr Stream stars to be present in this field. We have overlaid the same age+metallicity isochrones used in the ORPH-F1 field for identifying Sgr stream stars, but at a distance of 22 kpc based on Figure 4 of Belokurov et al. ( 2014 ). Among the stars consistent with these isochrones, seven are easily identified as a conspicuous clump in the PM diagram. These are plotted in blue circles in Figure 4 . For comparison, we also plot the Law Majewski ( 2010 ) N -body model particles that are within 15 in both right ascension and declination from the center of our HST field. The observed PM clump is very close to the model particles as expected, and similar to the findings for Sgr FIELD 4 in Sohn et al. ( 2015 ) at Λ ⊙ = 2511. Beyond these, no other significant excess has been found in the PM diagram. 3.3. ORPH-F3 Figure 5 shows the CMD and PM diagram for the ORPH-F3 field. If present, Orphan Stream stars should be at a distance of 24 kpc. We used the fiducial isochrones to identify Orphan Stream stars, and found four stars consistent with the isochrones within their photometric errors. However, all of these stars are far away from the PM prediction of Newberg et al. ( 2010 ). There is only one star (22) that is close enough to the model prediction to be considered a member of the Orphan Stream, but its color ( F 606 W − F 775 W ≃ 0.8) is too red for its brightness to be considered an Orphan Stream star. We conclude that there is no star associated with the Orphan Stream in this field. Given that this field is about 12 away in the orthogonal direction from the densest part of the Orphan Stream, and given that the Orphan Stream is a thin stellar stream, it is unsurprising to find no star in a small HST ACS field. Figure 5. Selection of stars associated with stellar streams in the ORPH-F3 field. The panels and symbols are similar to those in Figure 2 . The black solid curves in panel (a) are the same isochrones as used in Figure 2 , but for a distance of 24 kpc. The dashed green curve is a fiducial isochrone for a population with an age of 12 Gyr and a
Paper: 2017ApJ...845L...5I (fulltext)
then actually expect less active nuclear star formation and/or inflows in NGC 1097, which will lead to the deficiency of a geometrically thick torus in this LLAGN as compared to luminous AGNs such as NGC 1068 due to the deficiency of the energy source(s): this is an expected trend by the evolutionary torus models (Section 1 ). Another remarkable thing is that the envelope of the CND traced by the CO(3–2) emission and that traced by the NIR H 2 emission (Davies et al. 2009 ; Hicks et al. 2009 ) are almost identical in NGC 1097 at our available resolutions (Figure 3 (d)). This indicates that the outer radial extent of the molecular CND is unrelated to gas temperature, although the vertical height of the CND could differ (see below), which is consistent with hydrodynamic simulations of CND-scale gas distribution (e.g., Wada et al. 2009 ). However, we also found that their internal distributions differ significantly: while the two knots stand out in the CO(3–2) map, the NIR H 2 emission is clearly concentrated toward the nucleus, likely reflecting complex excitation conditions including fluorescence around the AGN. Thus, we should be careful to simply assume a linear conversion from NIR H 2 line luminosity to CO-based (Müller-Sánchez et al. 2006 ) when using that luminosity as a tracer of molecular mass. Gas dynamics is also key for studying the nature of the torus. Figures 4 (a) and (b) show the observed line-of-sight velocity and velocity dispersion maps of the CO(3–2) line in NGC 1097, respectively. It is evident that rotation (kinematic P.A. ≃ 118°) dominates the global gas dynamics, while streaming motions as well as high-velocity dispersions along the gas spirals/arc are detected. Figure 4. Intensity-weighted (a) mean velocity and (b) velocity dispersion maps of the CO(3–2) emission in the central 400 pc of NGC 1097. The stars mark the AGN location. The systemic velocity is 1254 km s −1 (thick line). Thin contours indicate 1080–1400 km s −1 in steps of 20 km s −1 for (a), and 10–40 km s −1 in steps of 10 km s −1 for (b), respectively. (c) and (d) Position–velocity diagrams (PVDs) of the CO(3–2) line along the adopted major axis (P.A. = 118°) and minor axis (P.A. = 208°) shown in the blue contours. The overlaid red contours indicate the PVDs produced by our dynamical model (MODEL-1). Both contours are plotted at 5, 10, ..., 40 σ , where 1 σ = 0.84 mJy beam −1 . Due to the inclined geometry of the concentric rings, no model component exists at offset ≳11 along the minor axis. To extract basic beam-deconvolved dynamical information, particularly the rotation velocity () and the velocity dispersion ( σ ), we fitted tilted rings to these observed maps with the 3D Barolo code (Di Teodoro Fraternali 2015 ). The main parameters are dynamical center, systemic velocity (), , σ , galactic inclination ( i ), and P.A., all of which can be varied for each ring. Our initial fits indicated a dynamical center as the AGN position, as 1254 km s −1 . Thus , σ ,
Paper: 2016A&A...592A..61L (acknowledgments)
We thank A. Strong and T. Porter for advice regarding the details of the GALPROP ISRF model used in this paper. We also thank A. M. Stutz, Y. L. Shirley, J. Steinacker, and E. Keto for helpful discussions and comments on the paper draft. The work of J.K. and S.E.R. was supported by the Deutsche Forschungsgemeinschaft priority program 1573 (“Physics of the Interstellar Medium”). H.L. and M.N. were funded by the Deutsches Zentrum für Luft- und Raumfahrt (DLR). PACS has been developed by a consortium of institutes led by MPE (Germany), including UVIE (Austria); KU Leuven, CSL, IMEC (Belgium); CEA, LAM (France); MPIA (Germany); INAF-IFSI/OAA/OAP/OAT, LENS, SISSA (Italy); IAC (Spain). This development has been supported by the funding agencies BMVIT (Austria), ESA-PRODEX (Belgium), CEA/CNES (France), DLR (Germany), ASI/INAF (Italy), and CICYT/MCYT (Spain).
Paper: 2017MNRAS.467.3934D (acknowledgments)
Acknowledgments This research was conducted by the Australian Research Council Centre of Excellence for All-sky Astrophysics (CAASTRO), through project number CE110001020. SB acknowledges funding support from the Australian Research Council through a Future Fellowship (FT140101166). GAMA is a joint European-Australasian project based around a spectroscopic campaign using the AAT. The GAMA input catalogue is based on data taken from the SDSS and the UKIRT Infrared Deep Sky Survey. Complementary imaging of the GAMA regions is being obtained by a number of independent survey programmes including GALEX MIS, VST KIDS, VISTA VIKING, WISE , Herschel -ATLAS, GMRT and ASKAP providing ultraviolet to radio coverage. GAMA is funded by the STFC (UK), the ARC (Australia), the AAO and the participating institutions. The GAMA web site is http://www.gama-survey.org/ .
Paper: 2015ApJ...812..107H (acknowledgments)
J.C.L. and K.M.H. acknowledge support from the INSU-CNRS national programs PCMI and PNP. V.W. is funded by the ERC Starting Grant (3DICE, grant agreement 336474). Acknowledgement is made to the Donors of the American Chemical Society Petroleum Research Fund for partial support of this research (PRF#53105-DN16 for FG summer support). The Rennes team acknowledges support from the Agence Nationale de la Recherche, contract ANR-11-BS04-024-CRESUSOL-01, the INSU-CNRS national program PCMI, the Institut National de Physique (INP CNRS), the Région Bretagne and the Université de Rennes 1. S.D.L.P. acknowledges financial support from the Institut Universitaire de France.
Paper: 2017ApJ...842L..22Z (fulltext)
is listed in Table 1 , and plotted in Figure 2 together with LAE LFs at z ∼ 5.7–7.3 from the literature. The LF shows dramatic evolution not only in normalization, but also in the shape. Figure 2. Top: the Ly α LFs of LAE surveys at z ∼ [5.7, 6.6, 7, 7.3]. The red filled circles show our results at z ∼ 7 from the LAGER-COSMOS field. References: Ly α LFs at z = 5.7 from Ouchi et al. ( 2008 ; blue filled squares), Konno et al. ( 2017 ; blue filled stars), and Santos et al. ( 2016 ; light-blue filled hourglass), at z = 6.6 from Ouchi et al. ( 2010 ; green filled squares), Konno et al. ( 2017 ; green filled stars), and Santos et al. ( 2016 ; dark-green filled hourglass), at z ∼ 7 from this work (red filled circles) and Ota et al. ( 2010 ; pink empty circles with plus), and at z = 7.3 from Shibuya et al. ( 2012 ; purple empty circles), and Konno et al. ( 2014 ; purple empty circles with cross). The colored regions in blue and green show the best Schechter fit with ±1 σ error in L * at z ∼ 5.7 and 6.6, respectively, from Konno et al. ( 2017 ). The solid and dashed lines in red (gray) show the best Schechter fits at z ∼ 6.9 with L range 42.65 ≤ log 10 L ≤ 43.25 (42.65 ≤ log 10 L ≤ 43.65) (see Table 2 ). A truncated Schechter function is also plotted (red solid line, see the text for details) to measure the bright-end LF bump. The purple line shows the best Schechter fit at z ∼ 7.3 from Konno et al. ( 2014 ). All these Schechter fits have a fixed faint-end slope of α = −2.5. Bottom: the histogram of Ly α luminosities of LAGER LAEs. The average completeness fraction for our LAEs interpolated from Section 2.2 is also plotted as a red solid line with red triangles. Table 1 Ly α Luminosity Function at z ∼ 7 from LAGER-COSMOS Field Note. a 1 σ Poisson errors from Gehrels ( 1986 ). The relatively fainter end of our LF (42.65 ≤ log 10 L Ly α ≤ 43.25) at z ∼ 7 can be fitted with a a single Schechter function in the form ofBecause our survey does not go far below L *, we fix α = −2.5 in our fitting, which is suggested from recent studies on Ly α LFs at z = 5.7 and z = 6.6 by Konno et al. ( 2017 ). Konno et al. ( 2017 ) report the largest narrowband surveys to date for LAEs at z = 5.7 and z = 6.6 taken with Subaru HSC on the 14–21 deg 2 sky, and their best-fit Schechter function analyses also include the smaller but much deeper narrowband surveys taken with Subaru Suprime-Cam (Ouchi et al. 2008 , 2010 ) at the corresponding redshifts. Considering that the Ly α LFs by Konno et al. ( 2017 ) are derived from the largest LAE samples at z = 5.7 and z = 6.6 with the widest luminosity range, and we have used consistent selection and analysis methods, we choose their Ly α LFs for comparison in the following sections. At z > 7, we choose the Ly α LF at z = 7.3 by Konno et al. ( 2014 ) for comparison. The best-fit Schechter function parameters of these LFs are listed in Table 2 and plotted in the upper panel of Figure 3 . We find a large drop in Φ of Ly α LF, compared
Paper: 2019ApJ...887..138W (acknowledgments)
D.A.D. and A.D.W. are grateful for funding from EPSRC via grant No. EP/N018117/1. C.R.S. is grateful for funding from the Carnegie Trust for the Universities of Scotland via research incentive grant No. RIG007788. A.D.W. gratefully acknowledges the support of NVIDIA Corporation with the donation of a Titan Xp GPU used in this research. Data used in this project are available at https://doi.org/10.5525/gla.researchdata.910 .
Paper: 2018ApJ...859...21A (acknowledgments)
M.A. and J.P.W. were supported by NSF and NASA grants AST-1208911 NNX15AC92G, respectively. M.A. also acknowledges support from NSF AST-1518332, NASA NNX15AC89G, NNX15AD95G/NEXSS, and the Center for Integrative Planetary Science at UC Berkeley. Leiden is supported by the European Union A-ERC grant 291141 CHEMPLAN, by the Netherlands Research School for Astronomy (NOVA), and by grant 614.001.352 from the Netherlands Organization for Scientific Research (NWO). A.M. and C.F.M. are supported by ESO Fellowships. M.R.H. is supported by NWO grant 614.001.352. M.T. has been supported by the DISCSIM project, grant agreement 341137 funded by the European Research Council under ERC-2013-ADG. This paper makes use of the following ALMA data: 2015.1.00222.S, 2016.1.01239.S, 2013.1.00226.S, 2013.1.01020.S. ALMA is a partnership of ESO (representing its member states), NSF (USA) and NINS (Japan), together with NRC (Canada), NSC and ASIAA (Taiwan), and KASI (Republic of Korea), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO and NAOJ. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. This work benefited from NASA’s Nexus for Exoplanet System Science (NExSS) research coordination network sponsored by NASA’s Science Mission Directorate.
Paper: 2017ApJ...849..113L (fulltext)
of Fe xxi 1354.09 Å. On the other hand, the shorter period cannot be detected in these detrended fluxes, which further confirms that the 3.1 minute oscillations can only be observed in the Doppler shift of Fe xxi 1354.09 Å. Figure 9. Normalized Fourier power spectra of the detrended time series from IRIS Fe xxi 1354.09 Å (a, b), GOES SXR fluxes (c, d), and AIA EUV light curves (e, f). A horizontal purple line in each panel indicates the 99.99% confidence level. As mentioned above, the Doppler shift oscillations within a period of ∼3.1 minutes ( P 1 ) can be interpreted as the standing kink oscillations in the hot flare loop. On the other hand, the longer period ( P 2 ) is different from the shorter one ( P 1 ), because it can be observed simultaneously in SXR and EUV fluxes, as well as in the time series of Doppler shift and line-integrated intensity of Fe xxi 1354.09 Å. Those facts imply the process of periodic energy release in this flare (Tian et al. 2015 ; Li et al. 2017b ). The longer period is most likely to be linked to the quasi-periodic pulsations during the solar flare, which seems to be caused by an MHD wave in the slow mode (Fang et al. 2015 ; Van Doorsselaere et al. 2016 ). Based on the formation temperature ( 11 MK) of Fe xxi 1354.09 Å, the local sound speed ( v s ≈ ) in the flare loop is estimated to be ∼500 km s −1 (e.g., Nakariakov Ofman 2001 ; Kumar et al. 2013 , 2015 ). While the deprojected length ( L ) and longer period ( P 2 ) of this flare loop have been identified as ∼81 Mm and ∼10 ± 4 minutes, then the phase speed of this flare loop ( v h = 2L/ P 2 ) can be estimated as about 200−450 km s −1 . This phase speed is slower than the local sound speed, implying it might be a slow MHD wave (Mandal et al. 2016 ). In other words, the longer period ( P 2 ) might be explained as the MHD wave in a slow mode. To compare the Doppler shift oscillations, we also use the sine function with a damping amplitude (Equation ( 3 )) to fit the time series of Doppler shift during peaks “5” and “6” (e.g., Nakariakov et al. 1999 ; Wang et al. 2002 ; Anfinogentov et al. 2015 ), because the Doppler shift oscillations are most pronounced during this time interval. Figure 8 (upper) shows the fitting result, the overlaid turquoise profile gives the best fitting, and it appears to match well with the Doppler velocities (black).Here, v 0 is the initial Doppler velocity, v m is the initial oscillatory amplitude, and P 1 , ψ , and τ are the shorter dominant period, initial phase, and decay time of the oscillations, respectively. 5. Conclusions Using observations from multi-instruments, i.e., IRIS , SDO /AIA, SDO /HMI, Hinode /XRT, and NoRH, we investigate the Doppler shift oscillations in a hot loop during a GOES M7.1 flare on 2014 October 27. The main results are summarized as follows. (I) The Doppler shift oscillations are detected in a hot flare line, i.e., Fe xxi 1354.09 Å from IRIS spectroscopic observations. The time series of Doppler shift
Paper: 2018A&A...612A..51B (fulltext)
noise random number generator within GILDAS . However, the data cube has undergone many transformations and the true noise may not be precisely Gaussian. Thus, we extract contiguous channels from signal-free regions of the cube which we use as noise. In fact, as any transformations were designed to preserve particularly the region where signal is present, this is a sort of worst-case noise. We first examine what we obtain from mock clouds with no velocity gradient, that is to say with v xy constant. Figure 13 shows the results for T max = 200 mK using both purely random noise and noise taken from signal-free but unoptimized regions of the cube. The gradients, for equivalent clouds and noise levels, are clearly more dispersed with the “real” cube noise. The dispersions, as measured by the full width at half power of the distribution (divided by 2.35 to give the equivalent for a Gaussian), are approximately 0.008 and 0.015 km s −1 pc −1 for the random and cube noise injection. T max = 200 mK corresponds roughly to the median signal in the cloud sample. Fig. 13 Histogram of gradients from noise. The black line shows the distribution of the gradients for purely random noise. The red line shows the distribution when noise is from signal-free regions of the cube. No velocity gradient has been injected here. We have also tested with higher and lower S/N levels. For T max = 100 mK, the dispersion increases to 0.014 and 0.029 km s −1 pc −1 for respectively random and cube noise. For T max = 400 mK, the dispersion decreases to 0.005 and 0.0085 km s −1 pc −1 for respectively random and cube noise. 5 Evaluating uncertainties on cloud rotation The same operations were done with the observed gradients, creating the same clouds but injecting the velocity gradient deduced from the observations for each cloud. Should there be a link between size or shape and the velocity gradient deduced from calculating the first moment, the link would be preserved in these tests. Figure 14 shows the distribution of the retrieved gradients. The clouds were created with the observed gradients and shapes and then noise was added and the gradients remeasured. This can be directly compared with Fig. 10 . The distribution in Fig. 14 is of course wider because the gradients from Fig. 10 have been injected and then noise added. The process was repeated with T max = 100 mK and the distribution is significantly wider but the prograde-retrograde asymmetry is nonetheless preserved. We are thus confident that the gradients retrieved reflect the true distribution of velocity gradients, although the distribution is likely broadened by the presence of the noise in the datacube. Fig. 14 Like Fig. 13 except that the observed gradients (see Fig. 10 ) have been injected, noise added, and the gradients recovered. Black line shows the distribution of the gradients for purely random noise. The red line shows the distribution when noise from signal-free regions of the cube. As in Fig. 13 , T
Paper: 2019ApJ...872...64K (fulltext)
) resolved the molecular line emission from above and below the disk, we could not detect the emission in the southwest of the disk, possibly due to the lower sensitivity and resolution. Figure 3. (a) Peak positions of the SiO v = 1 and 2 J = 1–0 maser emission (gray circle) and the SiO v = 4 J = 11–10 emission (colored circle). (b) Peak positions of the SiO v = 1 and 2 J = 1–0 maser emission (gray circle) and 29 SiO v = 2 J = 11–10 emission (colored cross). The color of the SiO v = 4 and 29 SiO v = 2 emission represents the line-of-sight velocity. The cross at (0, 0) corresponds to the position of Source I (Goddi et al. 2011 ). The 29 SiO v = 2 J = 11–10 emission is mainly found between the eastern and northern arms of the X-shape and also in the vertical structures from the disk, such as blueshifted emission along the eastern arm and high-velocity components near the disk midplane. It confirms that the 29 SiO v = 2 J = 11–10 emission traces the rotating outflow as other SiO line emission (e.g., Plambeck Wright 2016 ; Hirota et al. 2017 ; Ginsburg et al. 2018 ). On the other hand, the v = 4 J = 11–10 emission is only found in a line connecting the base of eastern and northern arms of 43 GHz SiO masers, which resembles the H 2 O emission in Ginsburg et al. ( 2018 ). The morphology and velocity gradient of this line suggests that it traces both the surface of the disk and the base of the outflow. 4.2. The Kinematics of <italic>v</italic> = 4 <italic>J</italic> = 11–10 Line Emission Figure 4 shows the position–velocity diagrams of the SiO v = 4 J = 11–10 and 29 SiO v = 2 J = 11–10 emissions along the axis with the position angle of 140° (Hirota et al. 2014 ). The spatial distribution and the velocity gradients along the major axis of the continuum emission indicate that the SiO v = 4 J = 11–10 emission is related to the rotating disk. However, the absence of the high-velocity components at the central position suggests that the high-frequency SiO line emissions are associated with a rotating ring structure rather than a filled disk. To find the parameters of the ring structure, we fit the observed rotation curve with a simulated rotation curve using a simple model similar to those of Hirota et al. ( 2014 ) and Plambeck Wright ( 2016 ). We assumed a thin, edge-on ring structure in Keplerian rotation around a central star. In each position, the gas with uniform temperature and density emit the optically thin emission. The thermal width of 3 km s −1 was assumed. We set the systemic velocity of 6 km s −1 , which is the central value of the observed velocity range. We simulated the rotation curve for a range of inner radii r in and outer radii r out with central masses ranges of 5 to 15 M ⊙ . As a result, we assess that r in = 12 au, r out = 26 au, and a central mass of 7 M ⊙ for the best parameter. Figure 4. (a) Position–velocity diagram of the SiO v = 4 J = 11–10 emission. The Keplerian velocity curves for 10 M ⊙ (gray) and 15 M ⊙ (black) are
Paper: 2017ApJ...835..262B (acknowledgments)
A.G.E. was supported by grant NNX10AT78G from NASA’s Goddard Space Flight Center. N.H.B. was supported by a STFC grant. We thank Gary Zank for useful comments.
Paper: 2016ApJ...821....5D (acknowledgments)
A.J.D. is supported by a Porat Fellowship at Stanford University. We thank Evan Kirby and Andrew Wetzel for stimulating science discussions, and for providing very useful comments on an earlier version of this manuscript. We also thank an anonymous referee for providing valuable comments that improved the quality of our paper. The simulations used were run using computational resources at SLAC; we gratefully acknowledge the support of the SLAC computational team. This research also used resources of the National Energy Research Scientific Computing Center, a DOE Office of Science User Facility supported by the Office of Science of the U.S. Department of Energy under Contract No. DE-AC02–05CH11231. A.J.D. thanks J.B.C. for his impeccable timing and invaluable insight.