adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016ApJ...828...84S (fulltext)
the case of PRD and the dotted line that of CRD. Other input parameters are . No background continuum opacity is used. 4.2. The Case of the <inline-formula> <tex-math> </tex-math> <inline-graphic xmlns:xlink="http://www.w3.org/1999/xlink" xlink:href="apjaa3333ieqn70.gif"></inline-graphic> </inline-formula> Transition For the case when J a = J b = 1, even in the absence of LLP, since the upper level is polarized, a finite amount of emergent polarization is generated unlike the case of J a = 1, J b = 0. For all the computations of this particular transition ( J a = J b = 1), we again consider the hypothetical case like that described for J a = 1, J b = 0. In this case, there are four density matrix elements to be determined, namely, , , , and , when polarizability of both levels are taken into account. Figure 2 shows the emergent ( I , Q / I ) at μ = 0.11 for the transition. The solid line in Figure 2 represents the emergent profiles obtained when both the effects of PRD and LLP are considered. In order to see the importance of both effects we have overplotted the ( I , Q / I ) profiles obtained when only the effects of PRD are considered with an unpolarized lower level (ULL—dashed line); when the effects of LLP is considered in the limit of CRD (dotted line); and the case where only the CRD effects are considered with ULL (dotted–dashed line). Figure 2 clearly shows that the LLP effects appear only in the emergent polarization and the intensity profiles remain unchanged whether or not LLP is taken into account. We see that the LLP effects in the emergent Q / I are significant mainly in the core (up to ∼2 Doppler widths, see inset in the lower panel of Figure 2 ), and in the wings the effects of PRD are dominant (compare solid and dashed lines). The enhancement in the emergent polarization at the line center when LLP is included is around 5%. Figure 2. Emergent intensity and polarization for μ = 0.11 computed using the two-stage approach for the case. The different line types correspond to: solid line—PRD + LLP, dotted line—CRD + LLP, dashed line—PRD + ULL, and the dotted–dashed line—CRD + ULL. The abbreviation ULL stands for unpolarized lower level. Other input parameters are . No background continuum opacity is used. The inset in the Q / I panel shows the Q / I profiles for a shorter frequency bandwidth for the sake of clarity. In the first stage of the two-stage approach, the SEEs are solved under the approximation of CRD, and thus all the transition rates that enter into the SEEs are frequency-integrated quantities. Therefore, all the redistribution effects are integrated away. The contributions to the frequency-integrated scattering probability come almost entirely from the Doppler core. In SEEs, we compare the individual transition rates for a given radiation field. For each individual transition, the contributions of the wing photons are insignificant compared with the core photons. Since it is only the core photons that are relevant
Paper: 2019ApJS..245...24W (fulltext)
analysis is beyond the scope of this work and will be discussed in Paper III. The results presented herein provide a statistically significant list of values and histogram distributions for reference and baselines in future work. Given the large statistics in this investigation, the beam/strahl fits for selection criteria Criteria UP are now the most comprehensive attempt to parameterize the beam/strahl electron velocity moments in the ambient solar wind. These results are useful for multiple modeling and simulation studies in addition to comparison with inaccessible regions like the intracluster medium. Further, these results will provide a statistical baseline for the Parker Solar Probe mission and the future Solar Orbiter and IMAP missions. Appendix A Definitions and Notation As in Paper I , in this appendix the symbols and notation used throughout will be defined. For all direction-dependent parameters we use the subscript j to represent the direction, where j = tot for the entire distribution, j = ∥ for the the parallel direction, and j = ⊥ for the perpendicular direction, where parallel/perpendicular is with respect to the quasi-static magnetic field vector, (nT). The generic subscript s is used to denote the particle species (e.g., electrons, protons, etc.) or the component of a single particle species (e.g., electron core). For the electron components, the subscript will be s = ec for the core, s = eh for the halo, s = eb for the beam/strahl, s = eff for the effective, and s = e for the total/entire population. Below are the symbol/parameter definitions: – one-variable statistics – X min ≡ minimum – X max ≡ maximum – – – X 25% ≡ lower quartile – X 75% ≡ upper quartile – σ ≡ standard deviation – σ 2 ≡ variance – fundamental parameters – ε o ≡ permittivity of free space – μ o ≡ permeability of free space – c ≡ speed of light in vacuum (km s −1 ) = – k B ≡ the Boltzmann constant (J K −1 ) – e ≡ the fundamental charge (C) – plasma parameters – n s ≡ the number density (cm −3 ) of species s – m s ≡ the mass (kg) of species s – Z s ≡ the charge state of species s – q s ≡ the charge (C) of species s = Z s e – T s , j ≡ the scalar temperature (eV) of the j th component of species s – the temperature ratio (N/A) of species s and s ′ of the j th component – the temperature anisotropy (N/A) of species s – V Ts , j ≡ the most probable thermal speed (km s −1 ) of a one-dimensional velocity distribution (see Equation 1 (c)) – ≡ the drift velocity (km s −1 ) of species s in the plasma bulk flow rest frame – C s ≡ the sound or ion-acoustic sound speed (km s −1 ) (see supplemental PDF file in Wilson et al. 2019b , for definitions) – V A ≡ the Alfvén speed (km s −1 ) (see supplemental PDF file in Wilson et al. 2019b , for definitions) – V f ≡ the fast mode speed (km s −1 ) (see supplemental PDF file in Wilson et al. 2019b , for definitions) – Ω cs ≡ the angular cyclotron frequency (rad s −1 ) (see Equation 1 (d)) – ω ps ≡ the angular plasma frequency
Paper: 2018ApJ...864...54D (fulltext)
the site S234. The data were observed with the Giant Metrewave Radio Telescope (GMRT) facility on 2016 December 8 (Proposal Code: 31_025; PI: L. K. Dewangan). The radio data reduction was carried out using the AIPS software, in a similar manner to that given in Mallick et al. ( 2012 , 2013 ). The synthesized beam size and the rms noise of the final 610 MHz continuum map are 56 × 46 and 120 μ Jy/beam, respectively. It is expected, specifically at low frequencies, that a sufficient amount of background emission will increase the antenna temperature while observing the source. Hence, we corrected the final GMRT map at 610 MHz for system temperature (see Omar et al. 2002 ; Mallick et al. 2012 , 2013 ; Baug et al. 2015 ). More details about the procedure adopted for correction of system temperature can be found in Mallick et al. ( 2012 , 2013 ). 2.2 Archival Data Sets The multi-wavelength data sets are also retrieved from different surveys (e.g., the NRAO VLA Sky Survey (NVSS; λ = 21 cm; resolution = 45′′; Condon et al. 1998 ), the Five College Radio Astronomy Observatory (FCRAO) 12 CO (1–0) and 13 CO (1–0) line data ( λ = 2.6, 2.7 mm; resolution ∼45′′; Heyer et al. 1998 ; Brunt 2004 ), the Herschel Infrared Galactic Plane Survey (Hi-GAL; λ = 70, 160, 250, 350, 500 μ m; resolutions = 58, 12′′, 18′′, 25′′, 37′′; Molinari et al. 2010 ), the Wide Field Infrared Survey Explorer ( WISE ; λ = 3.4, 4.6, 12, 22 μ m; resolutions = 61, 64, 65, 12′′; Wright et al. 2010 ), the Warm- Spitzer IRAC 3.6 and 4.5 μ m photometric data (GLIMPSE360; λ = 3.6, 4.5 μ m; resolution = 2′′; Whitney et al. 2011 ), the UKIRT NIR Galactic Plane Survey (GPS; λ = 1.25, 1.65, 2.2 μ m; resolution ∼08; Lawrence et al. 2007 ), the Two Micron All Sky Survey (2MASS; λ = 1.25, 1.65, 2.2 μ m; resolution = 25; Skrutskie et al. 2006 ), and the Isaac Newton Telescope (INT) Photometric H α Survey of the Northern Galactic Plane (IPHAS; λ = 0.6563 μ m; resolution ∼1′′; Drew et al. 2005 )). One can learn more details of these data sets and their analysis procedures in Dewangan et al. ( 2017e ). 3 Results 3.1 Extended Features in the Direction of l = 1717–1741 To explore the extended and embedded features toward l = 1717–1741, we have employed the submillimeter image at 250 μ m, a radio continuum map at 1.4 GHz (having an rms sensitivity of ∼0.45 mJy/beam; Condon et al. 1998 ), and 12 CO ( J = 1–0) line data (having a velocity resolution of 0.25 km s −1 and an rms sensitivity of 0.25 K; Heyer et al. 1996 , 1998 ). These data enable us to depict cold dust, ionized gas, and molecular gas in the sites. Figure 1 (a) displays the submillimeter image at 250 μ m overlaid with the 1.4 GHz continuum emission contours. Figure 1 (b) shows the 12 CO emission contour map at [−17.75, −1.0] km s −1 . We have labeled at least six objects (i.e., S237, S234, IRAS 05253+3504, IRAS 05231+3512, IRAS 05220+3455, and V582) in Figures 1 (a) and (b). The submillimeter image traces a spherical shell morphology of the site
Paper: 2020MNRAS.498.2415M (fulltext)
this is 0.3 ≲ f h ≲ 0.4 for halo masses M 12 ∼ (0.5−2) at redshift z ≳ 2 (Roca-Fàbrega et al. 2019 ), while observations of M 12 ∼ (1−10) haloes at z ≳ 0.1 suggest f h ∼ 0.2−0.3 (Singh et al. 2018 ). We define f h,0.3 = f h /0.3. Δ( R v ) is the ratio of the density at R v to the mean density in the halo, . For a singular isothermal sphere Δ( R v ) = 1/3, while for an NFW halo with concentration c = (5, 10, 20), Δ( R v ) ∼ (0.24, 0.19, 0.14), respectively. We write Δ 1/6 = Δ( R v )/(1/6). Assuming cosmological parameters Ω m = 0.3 and , we obtain for the stream density at R v where δ 100 = δ/100 and . For a hydrogen mass fraction of X = 0.76, this yields a hydrogen number density in the stream at R v of We hereafter write . The stream radius at R v can be constrained using the mass flux through the stream, where is the total mass accretion rate on to the halo virial radius, and f s is the fraction of baryonic accretion along the gas stream. Cosmological simulations suggest f s ∼ (0.2−0.5) with a typical value of f s = 1/3 (i.e. three significant streams; Danovich et al. 2012 ). In the Einstein de Sitter (EdS) regime (valid at z > 1), the accretion on to the virial radius is well approximated by (Dekel et al. 2013 ) with halo-to-halo variance encapsulated by the normalization s ∼ (0.5−2). Combining equations ( 15 ), ( 23 ), ( 25 ), and ( 26 ), we obtain with . Together with equation ( 16 ), we have Note that this is larger by a factor of ∼2 than the corresponding equation (68) in Padnos et al. ( 2018 ), which seems to be due to a typo, as it is consistent with their equation (67). Using equation ( 21 ) for δ, we see that and is independent of redshift. More massive haloes are thus fed by relatively narrower streams compared to their virial radii. We hereafter substitute . The hydrogen column density perpendicular to the stream axis is N H, 0 = n H, 0 R s . Using equations ( 24 ) and ( 27 ), this yields For haloes at z ∼ 2, this is large enough for the streams to be largely self-shielded against the UV background if the neutral fraction is x H i ≳ 10 −3 . In collisional ionization equilibrium (CIE), this is true for (e.g. Goerdt et al. 2010 ), i.e. Θ s ≲ 2. Of course, streams may still be susceptible to local sources of UV radiation, such as star formation in satellite galaxies located along the stream, or starbursts/AGN activity in the central galaxy. To summarize, given the halo mass and redshift, the precise temperatures of the stream and halo gas, Θ s and Θ h , respectively, determine the density contrast, δ. The stream Mach number, M b , is then interchangeable with η, the ratio of stream-to-virial velocity. The stream density is then determined by , the normalization of the hot halo gas density near R v . Finally, the ratio of stream radius to virial radius is set by , the normalization of the gas accretion rate along the stream. For η ∼ 1, , and , a stream at z = 2 has , (0.05−0.28), and (0.03−0.16) for δ = 30, 100, and 300, respectively.
Paper: 2018MNRAS.473.2493M (fulltext)
reduced to a negligible percentage (Magliocchetti etal. 2002 ; Mauch Sadler 2007 ). We then distinguished between AGN-powered galaxies and star-forming galaxies by means of equation ( 1 ) for z 1.8 and by fixing log 10 P cross ( z )23.5 [W Hz 1 sr 1 ] at higher redshifts. This procedure identifies 704 AGNs (corresponding to 33percent of the total radio population) and 1419 star-forming galaxies. With respect to our earlier work, these numbers imply an increase of a third of sources in both the AGN and the SF samples and allow us to draw conclusions on the properties of radio-selected sources with and without an FIR counterpart on much more solid grounds. Also, note that, due to the adopted selection criteria and thanks to the depth of the VLA-COSMOS survey, the AGN samples are complete with respect to radio selection at all redshifts z 3.5, i.e. the considered sample includes all radio-emitting AGNs selected at 1.4 GHz in the COSMOS field and endowed with a redshift determination z 3.5. 272 sources classified as AGNs and 901 sources classified as star-forming galaxies also show up in the Herschel maps. The redshift distribution of radio-selected AGNs is presented in Fig. 1 . The filled histogram shows the distribution of all AGNs, independent of their FIR emission, while the dashed (red) histogram represents that of those AGNs which also appear in the Herschel maps. The bottom panel highlights the ratio between these two quantities. Errorbars correspond to 1 Poissonian estimates. As already seen in Magliocchetti etal. ( 2014 ), also in the present case we have that the redshift distribution of radio-selected AGNs presents a marked peak at a redshift z 1. However, the better statistics of the present data allow us to identify a prominent tail which extends up to z 3.54 and which in Magliocchetti etal. ( 2014 ) only showed up as a secondary peak centred around z 2.5. There is no functional difference between the distribution of the parent AGN population and that which corresponds to AGNs with an FIR counterpart. One is the scaled version of the other, and the ratio between these two quantities is roughly constant and equal to the value of 0.5 throughout the whole redshift range probed by our data. Figure 1. Redshift distribution of COSMOS-VLA radio-selected AGNs with fluxes F 1.4GHz 0.06mJy. The AGN sample is complete up to z 3.5. The solid line represents all sources irrespective of their FIR emission, while the dashed line indicates those objects which are also FIR emitters. The bottom panel highlights the ratio between the two quantities. Errorbars correspond to 1 Poissonian estimates. Some more information on the sources under exam can be provided by investigations of the distribution of their radio luminosities, both in the presence and in absence of FIR emission. This is shown in Fig. 2 , whereby the left-hand plot shows the distribution of sources at all redshifts, while the one on the right-hand side that of sources divided into
Paper: 2019MNRAS.486.5614I (fulltext)
Fig. 1 were obtained with the 1.6-m telescope of the National Laboratory for Astrophysics (Brazil) on 2011 May 31 (Ivanova et al. 2018 ). The instrumental setup and data processing are presented in detail in Cavichia, Costa Maciel ( 2010 ) and Ivanova et al. ( 2016 ). Table 1 summarizes the characteristics of our observations: the date of observation (the mid-cycle time, UT date), the geocentric ( ∆ ) and heliocentric ( r ) distances, the phase angle of the comet ( α ), the position angle of the extended Sun–comet radius vector ( φ ), the filter (photometry) and grid (spectroscopy), the total exposure time during the night ( T exp ), and the mode of the observation. Table 1. Log of the observations of comet 29P/Schwassman–Wachmann 1. 3 MORPHOLOGY OF THE COMA As can be seen in Fig. 1 , the λ4277/39 Å filter transmits the radiation of CO + ions and the continuum, whereas the transmitted radiation flux through the λ5309/169 Å filter may truly be considered as the continuum, without CO + ions. Thus, to derive the pure ion image, we subtracted the composite image of the continuum that was obtained in the SED537 filter from the combined CO + image, controlling the particular light transmission of each filter and various continuum levels. We used values of reddening obtained from the spectral observations of comet 29P presented by Korsun et al. ( 2008 ) and Ivanova et al. ( 2016 , 2018 ). The images of the comet in the two filters were centred using the central isophotes, which was closest to the maximum brightness of the comet. The panels on the left of Fig. 2 present the intensity maps of the isolated CO + coma after subtraction of the continuum intensity from the observed intensity in the CO + band, and the panels on the right present the images of the dust coma of comet 29P. Figure 2. Isolated CO + (a, c, e, g) and dust (b, d, f, h) coma of comet 29P derived from observations on November 8, 2007, December 4, 2008, March 29, 2009 and December 26, 2009. All images are scaled to erg cm −2 s −1 Å −1 . The values near the white contours, corresponding to the column density at a cometocentric distance about 20 000 km, are presented in Table 2 . The sunward (⊙), north (N) and east (E) directions, and the cometary motion vector (V) are shown. The isolated CO + coma is clearly observed on all the dates of observation presented here (see Fig. 2 ): it is compact, asymmetric relative to the optocentre, and always elongated in the sunward direction, although this elongation does not coincide exactly with the direction to the Sun. The CO + coma observed on 2007 November 8 was only slightly elongated in the sunward direction, while on 2008 December 4 the comet had a very compact and strongly elongated CO + coma in the solar hemisphere to at least 60 000 km. On 2009 March 29, the comet had a small asymmetric coma with strong sunward elongation. In contrast to the CO + coma, the dust coma, formed by sunlight scattered from dust particles, was more extended and
Paper: 2016ApJ...823..128M (fulltext)
the positions on the SED of the observational wavelengths for various instruments of interest here: 250, 350, and 500 μ m for the SPIRE instrument on board Herschel ; 450 and 850 μ m for the SCUBA-2 instrument at the James Clerk Maxwell telescope; ∼1100 μ m for AzTEC at the LMT; 1400 μ m for SPT ; and 850, 1400, and 3000 μ m for ALMA . We have also highlighted the 5 σ detection limits for such instruments (attained in the deepest large-scale surveys undertaken so far or in the near future): S 250 ≈ 35 mJy, S 350 ≈ 40 mJy, and S 500 ≈ 50 mJy for SPIRE; S 450 ≈ 8 mJy and S 850 ≈ 2 mJy for SCUBA-2 ; S 1100 ≈ 1 mJy for AzTEC ; S 1400 ≈ 20 mJy for SPT ; and S 850 ≈ 0.42, S 1400 ≈ 0.11, and S 3000 ≈ 0.02 mJy for ALMA (500 hr on 100 arcmin 2 ). The accurate determination of the spectroscopic redshift for a large sample of dusty galaxies is a major problem. In order to probe the bright end of the SFR function at z ≳ 3, a strategy could be to preselect high-redshift sources using flux/color criteria from surveys conducted with Herschel or SCUBA-2 , and then perform a more accurate photometric (or even spectroscopic) redshift determination with observations on source by, e.g., AzTEC and ALMA . Commonly used preselection criteria (e.g., Dowell et al. 2014 ; Asboth et al. 2016 ) for high-redshift sources based on Herschel photometry involve looking for “350-peakers,” which are defined as sources with S 350 / S 250 ≳ 1 and S 500 / S 350 ≲ 1, or “500-risers,” which are defined as sources with S 350 / S 250 ≳ 1 and S 500 / S 350 ≳ 1. Actually, the uncertainties in the flux measurements make the distinction between peakers and risers quite loose; moreover, at z ≳ 4, the channel at 250 μ m refers to rest-frame wavelengths λ ≲ 50 μ m, where details of the SED due to different dust properties and a possible contribution from an AGN component can be relevant. Thus, here, we mainly focus on the color S 500 / S 350 . In Figure 9 (bottom panel), we plot its evolution with redshift and compare it with the measurements from ALMA /ALESS by Swinbank et al. ( 2014 ), finding reasonable agreement within the large uncertainties. Sources with S 500 / S 350 ≳ 0.8 are high-redshift z ≳ 3 candidates. However, the rather high limiting fluxes of Herschel cannot probe the SFR function much above z ∼ 5, since even sources with SFR ψ ≳ 10 3 M ⊙ yr −1 are too faint to be detected. To go much beyond z ≳ 4, the preselection based on the color S 850 / S 450 from SCUBA-2 photometry is much more efficient. In Figure 9 (bottom panel), we plot its evolution with redshift and compare it with the measurements from SCUBA-2 by Koprowski et al. ( 2016 ), again finding reasonable agreement. It can be seen that the color condition S 850 / S 450 ≳ 0.6 can be exploited to preselect candidate galaxies at z ≳ 4. In Figure 10 , we show our predictions for the differential counts at 500 and 850 μ m for the red sources preselected according to the color criteria discussed above. At 500 μ m, the
Paper: 2017ApJ...845...30M (fulltext)
in steps of 1′′ (a total field of view 15 × 119). The roll angle of the slits is –45° to the N–S direction on the Sun. Here, we focus on FUV spectra from both channels containing ∼1333–1337 Å (C ii lines), ∼1342–1344 Å (continuum), ∼1348–1351 Å (Fe xii line), ∼1353–1356 Å (Fe xxi and O i lines), and ∼1399–1406 Å (Si iv line), and on the NUV channel containing four spectral regions: 2791–2806 Å (Mg ii h and k lines), 2813–2816 Å, 2825–2828 Å, and 2831–2834 Å. IRIS also recorded SJIs for four passbands with a step of 2′′ and a total field of view 120 × 119: 1330 Å (C ii transition-region filter), 1400 Å (Si iv transition-region filter), 2796 Å (Mg ii k chromospheric filter), and 2832 Å (Mg ii wing photospheric imaging). 2.2. Solar Flare on 2014 March 29 An X1 flare observed on 2014 March 29 occurred in the active region NOAA 12017 with the GOES 1–8 Å flux starting to rise impulsively at 17:45 UT and peaking at 17:48 UT. IRIS started to observe the active region three hours before the flare, and the observations lasted a few minutes after the flare peak with the NUV exposure time of 8.00 s before 17:46:04.78 UT and 2.44 s after 17:46:13.98 UT, and a 75 s raster cadence (see e.g., Heinzel Kleint 2014 and Kleint et al. 2015 , 2016 ; Liu et al. 2015 ; Young et al. 2015 for detailed descriptions of IRIS observations of this flare). In contrast to the previous flare, the gradual phase was detected only at its beginning, and the flare loops are barely visible. Moreover, the slit rasters cover only the less pronounced parts of the loop system. We add this X-type flare because it was widely studied in the literature; however, our detailed analysis is focused mainly on IRIS spectra from the flare of 2015 June 22. To show the time lag between the flare ribbons and loops, light curves from a small part of the flare ribbons (red) and their corresponding loop tops (blue) in 1400 and 2796 Å are plotted in Figure 2 . It is seen that the loop tops peak about 3 minutes later than the ribbons in both 1400 and 2796 Å channels. Figure 2. Time lag between the flare ribbons and loops for 2014 March 29. Light curves from a small part of the flare ribbons (red lines) and their corresponding loop tops (blue lines) in 1400 and 2796 Å. Black lines represent the integral emission from the whole flare. Note that the sharp blue peak between 17:45 and 17:46 is due to leakage of saturated emission from the ribbon. 2.3. The Radiometric Calibration In our analysis, we used IRIS -calibrated data (Level 2 data), which are corrected for dark current, flat field, and geometry effects, and are wavelength calibrated. The radiometric calibration was performed in the following way. Intensities in the IRIS data are given in Data Number (DN) or counts (DN s −1 ). We perform radiometric calibration of the observed intensities to absolute units [erg s −1 cm −2 sr −1 Å −1 ]. To convert intensities to absolute units, we used the following calibration procedure:where I abs is the absolute