adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016ApJ...824..124E (fulltext)
complex structure within the overall galaxy mass functions. For instance, in increasingly higher halo mass regimes, we observe high-mass drop-offs that occur at higher galaxy mass. These drop-offs mark the natural boundary in the largest central galaxy mass for a given halo mass, as seen in Figure 5 (b). Figure 13. Breakdown of RESOLVE-B (top row), L r ,tot -limited RESOLVE-B (middle row), and ECO (bottom row) SMF and BMF into different group halo mass regimes. The four group halo mass regimes are M halo 10 11.4 M ⊙ (low-mass groups, often solo centrals, blue), M halo between 10 11.4 and 10 12 M ⊙ (intermediate-mass groups, green), M halo between 10 12 and 10 13.5 M ⊙ (large group, orange), and M halo > 10 13.5 M ⊙ (cluster, red, ECO only). The dark gray marks the overall SMF or BMF for each data set, and the gray hash-marked region denotes the incomplete regions for each mass function. The conditional mass functions are much more complex than the overall mass function, with pronounced bumps and dips. In both the RESOLVE-B and ECO intermediate (green) and large (orange) group halo mass regimes, we observe a peak in the SMF and BMF. For intermediate group halos, this peak occurs at ∼10 10.2 M ⊙ , right between the gas-richness threshold and bimodality mass scales. For large group halos, the peak occurs at ∼10 10.8 M ⊙ , above the bimodality mass scale and also near the knee of the overall mass functions. The peak in the large group halo regime appears to be causing the pronounced “bump” seen in the overall RESOLVE-B SMF and BMF. While the ECO data set also has this characteristic peak in its large group halo regime mass functions, the more gradual decline in the overall SMF and BMF for large galaxy masses seems to be due to the cluster galaxy population in ECO. The cluster population makes up 17% of galaxies with M star > 10 10.5 M ⊙ and 31% of galaxies with M star > 10 11 M ⊙ for ECO. The RESOLVE-B “bump,” meanwhile, may be emphasized by the lack of a cluster galaxy population and the overabundance of halos in this large group halo regime, as shown in Figure 5 (a). It is clear that the low-mass slopes of the SMF and BMF in these different halo mass regimes are very different. In the cluster halo mass regime, the rise is quite smooth, although the slope appears to flatten out at lower masses, at least in part due to incompleteness in galaxy counts around the clusters caused by high fiber collision rates and by missing ultradiffuse galaxies like those recently found in the Coma and Virgo clusters (Mihos et al. 2015 ; van Dokkum et al. 2015 ). We note that Yamanoi et al. ( 2012 ) found a fairly flat slope for the Coma cluster LF down to galaxy magnitudes of M R = −14.0, with a steep upturn for galaxies below our luminosity limit. Also, not all previous studies of cluster LFs have found steep slopes, and many cluster LF studies have relied on using statistical counts to remove background galaxies (e.g., Dressler 1978 and Goto et al. 2002 ).
Paper: 2019MNRAS.483.2592H (fulltext)
effects of binary disruption and binary segregation resulting in a significant delay of complete mixing compared to the inert binary case. In order to further describe the overall effect of the different dynamical processes affecting the survival and evolution of binaries on the FG and SG binary mixing, we present in Fig. 3 the time evolution of the cumulative radial distribution of FG and SG binaries for four different models with different values of the initial hardness parameter of the binary population. Figure 3. Cumulative radial distribution of FG (red lines) and SG (blue lines) binaries. From left to right, each panel shows the evolution of radial distribution. Three different models with x g,0 = 3, 3-20, 50, and 800 are presented (top to bottom). The cumulative radial distribution of inert binaries are shown as dashed lines for comparison. The last panel for the model with x g,0 = 800, the distribution at t ∼ 40τ rh, 0 is shown to illustrate a complete spatial mixing of very compact binaries. The cumulative radial distribution of FG and SG inert binaries in the MPr5f1m2 are also shown in order to illustrate the spatial mixing of stars with the same mass of binaries but not affected by other processes such as ionization and ejection due to the recoil velocity acquired during interactions with other stars. At t = 0, FG and SG binaries in all models have the same cumulative distribution. For the MPr5f1x3 model, the rapid disruption of SG binaries significantly depletes the inner binary population as illustrated by the shift of the SG binary cumulative radial distribution towards the outer regions. For the MPr5f1x3-20 model, the evolution of the distribution of FG binaries is similar to that for the MPr5f1x3 model. However, the presence of harder binaries implies that the effects of ionization are not as strong as in the previous case and the evolution of the SG cumulative radial distribution is more similar to that of inert binaries. It is interesting to note the development of an approximately flat portion in the cumulative radial distribution of SG binaries in the MPr5f1x3 and x3-20 simulations. The formation of this shell depleted in SG binaries is due to the combined effect of binary segregation and disruption; binaries are rapidly disrupted in the inner regions but as the segregation time-scale increases with the distance from the cluster’s centre. Therefore, there is no replenishment of binaries beyond the gap because the segregation time-scale is longer than the cluster’s age (see also the fig. 7 in Paper I ). The development of this gap is also responsible for the rapid variation in R h,SG and R h,FG / R h,SG around t ∼ 23τ rh, 0 in Figs 1 and 2 for the MPr5f1x3-20. For the MPr5f1x800 model, the cumulative distributions for FG and SG binaries are, in general, very similar to those in the inert binary model. Eventually, the spatial mixing of FG and SG binaries is achieved at t ∼ 40τ rh, 0 . Fig. 4 shows the time evolution of the
Paper: 2019ApJ...882L..27D (fulltext)
we sample 10,000 parameter combinations and generate one random light curve with the time sampling of the observed Keck data (∼30 nights of historical data and four nights of 2019 data) for each parameter set. We present the complementary cumulative distribution function (CCDF, or 1—CDF) of these 10,000 light curves in Figure 5 in the form of the median CDF and 1, 2, and 3 σ credible intervals. Figure 5. Top panel: comparison of the complementary cumulative distribution function (CCDF) of the observed data (historical data and 2019 data; black) and the median CCDF (dashed blue line) and the 1, 2, and 3 σ contours calculated from 10,000 simulations. The simulations were drawn from the posterior in Witzel et al. ( 2018 ). The dashed section of the observed CCDF represents flux densities that occurred only during the brightest flux excursion on 2019 May 13. These simulations show that if we repeated the entire experiment with the time sampling of 30 historical nights of Keck observations 10,000 times, then the probability of observing a single night with flux levels as high as 6 mJy is less than 0.3%. Bottom panel: the same as for the top panel but contours determined from simulations based only on the time sampling of the four nights in 2019. Because three of the four nights have elevated Sgr A* flux levels, if an experiment with four nights of observations were repeated 10,000 times, the probability of observing Sgr A* flux levels similar to the nights in 2019 would be less than 0.05%. Based on simulations drawn from the Sgr A* flux model, we can compute the probability of observing the light curves presented here. The simulations show that if we repeated our experiment 10,000 times with the time sampling of the 30 nights of historical Keck observations (from 2005 to 2013) of Sgr A*, and including the 2019 nights, there is a 0.3% probability that we observe a single night with flux levels as high as seen on 2019 May 13. These simulations are the most consistent with our current observations because they have very similar noise properties, observed duration, and timing. If we also include VLT nights for a total of over 130 nights, then the probability is less than 1.5% to observe flux values higher than 6.18 mJy on a single night. We can also consider the probability of randomly drawing 4 light curves similar to the observations in 2019. This probability is considerably lower at less than 0.05%. We also note that we have only observed the decay of the light curve on 2019 May 13, which suggests the actual maximum was likely even higher (Figure 3 ). Here we examine two possibilities for explaining the very unusual brightness and variation of Sgr A* observed this year: (1) the statistical models need to be changed or updated, and (2) there is a physical change in the accretion activity of Sgr A*. Based on the statistical model of Witzel et al. ( 2018 ), with the four nights of observations, there is a probability of less than 0.05% to
Paper: 2016A&A...592L...7T (fulltext)
1 . While ID11 and notably also A4.1 fall in the transition region between HII galaxies and giant HII regions, the rest of the lensed systems are among the most luminous star-forming systems and occupy a region of the diagram corresponding to the most luminous HII galaxies. The Lynx arc result from Fosbury et al. (2003) is shown with two different values of velocity dispersion; the lower value corresponds to the semi-forbidden lines and the higher value corresponds to the fit to the permitted CIV 1550 Å line. There is more scatter in the lensed systems, but on the whole the results tend to support the lensing model estimates of the magnification apart from the Lynx arc result when using semi-forbidden lines. There is a crucial general point that we need to make regarding this exercise: The luminosity in these young bursts of star formation evolves quickly in timescales of few megayears, which moves their position away from the correlation to lower luminosities and smaller equivalent widths (EW) of H or H ( Melnick et al. 2000 ; Bordalo Telles 2011 ). This fact may be related to the scatter of the Richard et al. (2011) points below the relation (Fig. 1 ). It is therefore highly recommended to obtain and publish the EW of H or H of star-forming systems to allow for an evolution correction to be performed. 4. Conclusions Regarding the intrinsic properties of ID11, Fig. 1 implies that it is either a luminous giant HII region, such as 30 Doradus in the Large Magellanic Cloud or those found in spiral galaxies like M101, or a low luminosity HII galaxy. Given that the deep HST images of ID11 show no indication of an underlying galaxy, we confirm the Vanzella et al. (2016) conclusion that ID11 is a low luminosity HII galaxy at . The relation spans more than three orders of magnitude in luminosity, which allows us to observe a significant number of objects at large redshifts without recourse to gravitational telescopes; hence this accounts for the power of the relation as a cosmological probe ( Plionis et al. 2011 ; Terlevich et al. 2015 ; Chávez et al. 2016 ). But, of course, only the most luminous HII galaxies can be observed at higher redshifts with the current generation of 810 m telescopes, so gravitationally amplified objects such as ID11 provide us with the unique chance to verify that at high redshift the correlation holds also for low luminosity and possibly low metallicity HII galaxies. We used an independent method, which is based on the standard candle provided by the relation that is valid for HII galaxies and giant HII regions, to measure the amplification affecting the star-forming system ID11 by the Abell cluster S0163 as ; this coincides with the value of 17 obtained by Vanzella et al. (2016) from the strong lensing model of Caminha et al. (2015) . Our result suggests that we can use giant HII regions in high- lensed galaxies as instruments for studying gravitational lenses and that the combination of the standard candle plus detailed
Paper: 2016ApJ...817..114Z (acknowledgments)
We acknowledge the support of the staff of the Li-Jiang 2.4 m telescope (LJT), Xing-Long 2.16 m telescope, and Tsinghua-NAOC 0.8 m telescope (TNT). Funding for the LJT has been provided by the Chinese Academy of Sciences (CAS) and the People’s Government of Yunnan Province. The TNT is owned by Tsinghua University and jointly operated by the National Astronomical Observatory (NAOC) of the CAS. Financial support for this work has been provided by the National Science Foundation of China (NSFC, grants 11403096, 11178003, 11325313, 11133006, 11361140347, 11203034, 11473063, 11322327, 11103072, 11203078, 11303085, 11203070 and 11103078); the Major State Basic Research Development Program (2013CB834903); the Strategic Priority Research Program “The Emergence of Cosmological Structures” of the CAS (grant No. XDB09000000); the Key Research Program of the CAS (Grant NO. KJZD-EW-M06); the Western Light Youth Project; the Youth Innovation Promotion Association of the CAS; the Open Project Program of the Key Laboratory of Optical Astronomy, NAOC, CAS; and the key Laboratory for Research in Galaxies and Cosmology of the CAS.
Paper: 2017MNRAS.470.3103K (acknowledgments)
Acknowledgments DK and DP thank the Institut Lagrange de Paris, a LABEX funded by the ANR (under reference ANR-10-LABX-63) within the Investissements d'Avenir programme under reference ANR-11-IDEX-0004-02, where this work was started. AL acknowledges the NSF grant AST 1212096 and Center for Magnetic Self Organization (CMSO). He also acknowledges Institut D'Astrophysique de Paris and Institut Lagrange de Paris for hospitality during his visit.
Paper: 2016ApJ...820..134H (fulltext)
also shows a U-shaped cavity wall encompassing the E1 and E2 lobes, extending to ∼17 from the source. The emission is from the swept-up material in the dense core, which has a radius of 1500 au (∼17). Unlike the observation, however, no emission is seen extending over ∼17 from the source to ∼4′′ to the northeast. In our model, that region is within the AGB halo. Thus, in order to produce detectable emission there, the AGB halo could be denser than currently assumed in our model. Figure 6 shows the comparison of HV CO J = 3–2 emissions between the observation and the model. Both panels are also plotted over the HST images of CRL 618. In the observation (Figure 6 (a)), HV (−183.2 to −114.1 km s −1 ) emission is observed at the tips of lobes E1, E3, E4, and E5, but not observed at the tip of lobe E2. In the model, HV emission, with a similar velocity range to that in the observation, is predicted to be seen at the tips of all the lobes. The velocity of the emission increases with distance, similar to that seen in the observation. Notice that the emission intensity in our model is about half of that observed, and thus more bullet mass may be needed to produce more HV CO emission. The reason for no HV emission being detected at the tip of and along the southern edge of lobe E2 is unknown. Figure 6. (a) Observed HV CO J = 3–2 emission (contours) plotted on the H α image of CRL 618 (from Figure 2 of Lee et al. 2013a ). The velocity ranges of the contours are −183.2 to −160.6 km s −1 (blue), −157.8 to −132.5 km s −1 (cyan), and −132.5 to −114.1 km s −1 (green). The contours start at 1 Jy beam −1 km s −1 , with a step of 2 Jy beam −1 km s −1 . The synthesized beam is ∼053 × 035 at a position angle of ∼83°. (b) Simulated HV CO J = 3–2 emission (contours) in Model 2, with similar velocity ranges to those in (a). HV CO emissions are seen at the tips of the E1 to E5 lobes. The contours start at 0.5 Jy beam −1 km s −1 , with a step of 1 Jy beam −1 km s −1 . The synthesized beam is the same as that in (a). Figures 7 (a) and (b) show the PV diagrams of the observed outflows in CO, cut along the axis of lobes E1 and E2. Figures 7 (c) and (d) show the same PV diagrams but calculated from the model. In these diagrams, lobes E4 and E5 can also been seen. In Figures 7 (c) and (d), the bullets are seen at the tips of the outflows and thus are traced by the fast molecular outflows. The diagrams also show that the velocities of the molecular outflows decrease linearly with the distance to the central source, as seen in the observations. The slope of the PV structure is roughly equal to , where p is the projected distance from the center and t is the time after ejection. Therefore, different slopes are seen for the outflows in different ejection episodes. Notice that since no molecular emission is detected for lobe E2, no PV structure is detected for lobe E2 in the observation, as shown in Figure 7 (b). Figure 7. (a and b) PV diagrams of CO J = 3–2 emission of CRL
Paper: 2018ApJ...863..130W (acknowledgments)
This work was supported by NSF grants AST-1211995, AST-1616598, and AST-1616710. A.W. gratefully acknowledges the assistance of Vincenzo Pota in prototyping the modeling code. A.J.R. is a Research Corporation for Science Advancement Cottrell Scholar. A.V. is supported by the NSF Graduate Research Fellowship Program. D.F. thanks the ARC for financial support. J.S. acknowledges support from a Packard Fellowship. S.B. acknowledges the support of the AAO PhD Topup Scholarship. We thank the anonymous referee for useful suggestions. Some of the data presented herein were obtained at the W.M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California, and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation. This research has made use of NASA’s Astrophysics Data System and the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. Funding for the DEEP2/DEIMOS pipelines (Cooper et al. 2012 ; Newman et al. 2013c ) has been provided by NSF grant AST-0071048. The DEIMOS spectrograph (Faber et al. 2003 ) was funded by grants from CARA (Keck Observatory) and UCO/Lick Observatory, a NSF Facilities and Infrastructure grant (ARI92-14621), the Center for Particle Astrophysics, and by gifts from Sun Microsystems and the Quantum Corporation. This research made use of Astropy, a community-developed core Python package for Astronomy (Astropy Collaboration et al. 2013 ). We acknowledge the use of other open source Python packages, including Numpy (Walt et al. 2011 ), Scipy (Jones et al. 2001 ), Matplotlib (Hunter 2007 ), Pandas (McKinney 2010 ), IPython (Pérez Granger 2007 ), and Scikit-learn (Pedregosa et al. 2011 ). The authors wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community. We are most fortunate to have the opportunity to conduct observations from this mountain. Facility: Keck:II (DEIMOS).
Paper: 2016A&A...588A..23A (fulltext)
other hand, in more extended BCDs the observed CO emission corresponds only to the most luminous regions of the starbursts, so it may be considered as a lower limit to the total CO emission. For all the observed galaxies our CO observations cover at least 50% of the projected size of starburst (, see Sect. 3), while for three objects this fraction, , is 100%. Table 3 Multiwavelength properties of galaxies from the literature. Notes. Columns: (2) radius at the mag arcsec isophote; (3) radius of the star-forming region from 2D surface photometry ( Amorín et al. 2009 ); (4) inclination; (5) H i gas mass; (6) logarithm of H i gas surface density; (7) stellar mass; (8) logarithm of star-formation rate density; (9) gas-phase metallicity; (10) logarithm of H surface density assuming a Galactic CO-to-H conversion factor. Reference. For metallicity: References for : (a) Shi et al. (2006) ; (b) Cairós et al. (2007) ; (c) García-Lorenzo et al. (2008) ; (d) López-Sánchez Esteban (2010) ; (e) Zhao et al. (2010) ; (f) Pérez-Montero Díaz (2003) ; (g) Gil de Paz et al. (2003) . (h) Frayer et al. (1998) ; (i) Leroy et al. (2005) ; (j) Barone et al. (2000) ; (k) Albrecht et al. (2004) ; (l) Leon et al. (1998) ; (m) Verdes-Montenegro et al. (1998) ; (n) Taylor et al. (1998) ; (o) Tacconi Young (1987) ; (p) Wei et al. (2010) . A wobbler switching mode was used at a frequency of 0.5 Hz with a beam throw up to 240 in azimuth. Two independent SIS receivers at each frequency (A100, B100 at 3 mm and A230, B230 at 1.3mm) were used to observe both polarizations of the CO 10 and 21 transitions simultaneously. The receivers were tuned in single side band, using only the lower band (LSB), which improves the calibration procedure and avoids contamination by other spectral lines in the image band, especially in observations of the calibration sources. The SIS receivers were connected to two filter-banks with resolutions of 1 MHz and 4 MHz at 115 GHz and 230 GHz, respectively, which yields velocity resolutions, , of 2.6 km s and 5.2 km s, respectively. These backends have 512 channels each, covering a total instantaneous velocity range of 1300 and 2600 km s. Pointing and focus were checked before each integration on continuum sources and planets. The pointing accuracy was better than 2.5 on average. Calibrations were done using the chopper wheel method, by observing strong radio sources (Orion A, CW Leo, and W51d). System temperatures at both frequencies varied between 290–470 K, the higher values being due to bad weather conditions. Thus, we obtained a range for the antenna temperature sensitivity of –15 mK. The CO spectra was reduced using CLASS 2 ( Buisson et al. 1997 ). To increase the signal-to-noise ratio (S/N), averaged spectra for each pointing were obtained by weighting multiple scans by a factor of , where is the integration time and is the system temperature. In all cases, baselines of zeroth order were subtracted. The spectra was then scaled to a main-beam