adsabs/WIESP2022-NER huggingface dataset:
Paper: 2018ApJ...856...66M (fulltext)
and wiggles a few times before ending at a hotspot 15′′ from the core. No corresponding features are detected in the X-ray band. Similarly, there are no clear knot associations in the HST image. 1116+128 (4C+12.39). The extended radio emission consists of a single knot at 25 northwest of the core. The knot is not detected in the Chandra data. No features are apparent in the HST image at the radio knot. 1303−827 (PKS 1302−82). The extended radio emission consists of a single knot at 65 southeast of the core and a knot pair about 16′′ north-northwest of the core. The X-ray jet searched was set to the southeast to include the closest knot. No knots are detected by Chandra . In the HST image, there is faint, resolved flux from the closer of the north-northwest knot pair, shown in a inset in Figure 6 . 1502+106 (B2 0954+25A). The VLBI jet is oriented at a PA of 120°, while the kpc-scale jet is at a PA of 160° (see Table 10 ) and is marginally detected. The VLA image of the jet published by Cooper et al. ( 2007 ) shows a continuous jet all the way to the termination, about 9′′ from the core. Lister et al. ( 2013 ) found a maximum apparent speed of VLBI knots of 17.5 c . The HST image shows no clear detection of any part of the kpc-scale jet. 1622−297 (PKS B1622−297). The kpc radio jet is weak and not detected with Chandra . There is a pc-scale jet at a PA of −70°, with a maximum apparent speed of 18.6 c (Lister et al. 2013 ). There are no apparent features in the HST image that are associated with radio or X-ray emission. 1823+568 (4C+56.27). There is a pc-scale jet at a PA of −160° with a maximum apparent speed of 26.17 c (Lister et al. 2013 ). The kpc-scale jet is first oriented almost due south, but bends about 1′′ from the core to the east, where it is clearly detected in the Chandra image at 25 from the core. There are no apparent features in the HST image that are associated with radio or X-ray emission, although the southern portion lies along a diffraction spike. 2201+315 (4C+31.63). The X-ray image was first published by Hogan et al. ( 2011 ). This quasar has VLBI components with apparent velocities up to 8.3 c along a PA of −145° (Lister et al. 2013 ). In the VLA image, the jet is −37′′ long at a PA of −110°, terminating in a hotspot (Cooper et al. 2007 ). To the northeast, there is a lobe about 45′′ from the core. The radio jet is detected in the Chandra data only out to about 4′′ from the core. The quasar was not part of our observing program, so we do not have an HST image of it. 2230+114 (CTA 102, 4C+11.69). There is a pc-scale jet at a PA of 160° with a maximum apparent speed of 8.6 c (Lister et al. 2013 ). There are radio-emitting knots to the northwest and southeast sides of the core but no clear association with X-ray emission. There are no apparent features in the HST image that are associated with the radio knots. 4. Discussion 4.1. Detection Statistics We detected 9 X-ray emitting jets among the 17 sources that complete our
Paper: 2018A&A...610A..53M (fulltext)
form the Herschel maps, into radio signal at this frequency using the Bonzini et al. (2015) relation after scaling for the magnification factor does not discard the possibility that an AGN contribution might be responsible for the remaining emission. 3 <italic>Chandra</italic> X-ray observations The galaxies SDP.9 and SDP.11 were observed by Chandra in Cycle 16, with the ACIS-S3 CCD at the aimpoint, for a total exposure time of 55.3 ks (on January 10, 2015) and 19.8 ks (on February 28 and March 23, 2015), respectively. Very faint mode was used for the event telemetry format, and ASCA grade 0, 2, 3, 4, and 6 events were used in the analysis, which has been carried out using the ciao v4.7 software. Both sources are clearly detected, with 14 counts each in the 0.5–7 keV band within a circular region of 2 ′′ radius centered on the positions reported in Sect. 2 , and 8 counts in the observed 2–7 keV band. The 0.5–7 keV Chandra images of SDP.9 and SDP.11, centered on these positions, are shown in Fig. 1 . Remarkably, the 0.5 arcsec Chandra on axis resolution is able to resolve the X-ray emission from the lensed components. The X-ray emission better matches in position the SMA data than the HST images. Fig. 1 HST/WFC3 lens-subtracted F160W images (brightness scale in μ Jy arcsec −2 , left panel ) and Chandra 0.5/7 keV images ( right panel ) of SDP.9 and SDP.11, overlaid with signal-to-noise ratio (S/N) contours at 880 μ m from the SMA (in steps of 3, 6, and 9, Bussmann et al. 2013 ; Negrello et al. 2014 ). The lownumber of detected photons significantly limits the spectral constraints that can be obtained from the X-ray analysis. To provide basic information for both sources, however, we fit the spectra extracted from the Chandra data using simplemodels (within xspec v.12.9, Arnaud 1996 ) and the Cash statistics ( Cash 1979 ), which is appropriate for low-count spectra. Fitting the spectrum of SDP.9 with a power law provides Γ = 1.1 ± 1.0, which is suggestive of obscuration. When we fixed Γ to 1.8, as is typically observed in AGN and quasars (e.g., Piconcelli et al. 2005 ), we obtained a column density of ≈ 5.2 × 10 22 cm −2 at the source redshift (3 σ upper limit of ≈ 3.3 × 10 23 cm −2 ). The observed 2–10 keV flux is ≈ 3.6 × 10 −15 erg cm −2 s −1 , and the derived apparent rest-frame 2–10 keV luminosity (corrected for absorption but not for magnification) is ≈ 3.7 × 10 43 erg s −1 . The faint flux of SDP.9 explains the non-detection of the source (with an upper limit of ≈ 10 −14 erg cm −2 s −1 in the 0.5–8 keV band) in the XMM-ATLAS survey ( Ranalli et al. 2015 ). This survey covers a 6 deg 2 area with an average exposure of 3 ks. The non-detection of the AGN in XMM-Newton suggests that it is obscured, with an estimated column density of at least 10 23 cm −2 , as observed in other bright submm galaxies ( Johnson et al. 2013 ). For SDP.11, we combined the photons of the two observations into a single spectrum. As for SDP.9, a power-law model
Paper: 2016MNRAS.462.1029S (fulltext)
values at the reference epoch at which the two timing models of Table 4 are defined, MJD 55655. The fit to the position at the median MJD of each aeon (the finely dotted orange line) returns μ α = 5.63 ± 0.10 mas yr − 1 and μ δ = 2.34 ± 0.28 mas yr − 1 while the dashed lilac line represents the values obtained from the BTX model, shown in Table 4 . Using an estimated distance of ≃1040 pc (from the NE2001 model of free-electron distribution in the Galaxy; Cordes Lazio 2003 ) and a total proper motion , a 2D transverse velocity of ν t = 30 ± 9 km s − 1 was calculated. This assumes an uncertainty of 30 per cent 4 in the DM-derived distance mentioned above. The measurement is in agreement with the value of 30 ± 20 km s − 1 measured by Stappers et al. ( 1998 ) and represents a twofold increase in precision, even though the uncertainty of the DM-derived distance is assumed to be much greater. It should be noted that this is significantly lower than the average value of 93 ± 13 km s − 1 reported in Desvignes et al. ( 2016 ) for the transverse velocities of binary MSPs. However, it agrees well with the value of 56 ± 3 km s − 1 reported for the binary MSPs with distance measurements from parallaxes. The proper motion values obtained from the BTX model appear to be inconsistent with those obtained from fitting to position measurements for every aeon using the respective ELL1 models. This is because the proper motion terms and the OFDs are strongly covariant and therefore the uncertainties in the values obtained from the BTX model are heavily underestimated, reinforcing the need for cautious usage of this model. 3.2 DM variations Since the DM is a measure of the density of the ionized interstellar medium (IISM) along the line of sight to the pulsar, both the motion of the pulsar and the dynamical evolution of the IISM affect this value. While it is possible to obtain the DM from timing, ‘JUMPS’ or instrumental offsets introduced to align the ToAs from the different backends are fully covariant with the DM and prevent an accurate measurement directly from the data set presented above. Therefore, a DM value of 20.7299(17) cm − 3 pc is adopted from the LOFAR measurements of Kondratiev et al. ( 2016 ). When simultaneous dual (or multi-) frequency observations are available, it is, however, possible to accurately estimate the variation in the DM. The WSRT PuMa-II backend provides observations centred at 345 and 1380 MHz, with a cadence of roughly three weeks. Observations between the two frequencies are sequential, which are separated by, at most, a few days and available for the MJD range ∼54600–56800. Since low-frequency observations are more sensitive to the DM variations, these are utilized to measure them instead of the two frequency-band observations of the PSRIX backend. To measure DM variations, the PuMa-II ToAs were fitted for DM using the ELL1 model presented in Table 4 . The ToAs were then split into 100 d long intervals, to ensure that enough
Paper: 2017MNRAS.465..337P (fulltext)
shorter than the radiative length-scale, over which substantial radiative energy losses would occur. These results are independent of black hole mass. Figure 5. Figures (a–f) show the distance travelled by the jet before reaching equipartition for a range of values of the reconnection parameters A rec and d , from ( 97 ) and fractional Eddington jet power f Edd . We stop integration of the jet evolution equations at a distance of 10 8 r s if the jet plasma has not yet reached equipartition, since radiative energy losses at larger distances than this are insignificant. The jet parameters used to obtain these results are shown under Model C in Table 2 and the radiative losses corresponding to these models are shown in Fig. 4 . We find that the dimensionless distance to equipartition, X equi = z equi / r s , depends weakly on the fractional Eddington accretion rate, f Edd . The distance to equipartition decreases slowly with increasing f Edd . These results are independent of black hole mass, M . Figure 6. The desirable region of parameter space for jets with different fractional Eddington luminosities and reconnection parameters from Figs 4 and 5 . Here, ‘desirable’ means that the jet reaches equipartition at a distance, z equi , between 100 and 10 8 r s , whilst not radiating more than 95 per cent of its initial energy in reaching this point. The unfavoured regions of parameter space which do not fulfill these criteria are coloured in black. The jet parameters used to obtain these results are shown under Model C in Table 2 . Due to more severe radiative losses at higher fractional Eddington jet powers ( f Edd ), the desirable region becomes more restrictive for higher values of f Edd . This means that lower power jets are able to reach equipartition at smaller distances whilst retaining at least 5 per cent of their initial energy. For jet powers close to the Eddington luminosity ( f Edd > 0.1), the strong radiative losses essentially prohibit the jet from coming into equipartition at a distance smaller than ∼10 4 r s , since otherwise, the jet radiates away more than 95 per cent of its initial power in the compact base region with high magnetic field strengths. This is consistent with our earlier results in Fig. 3 , in which we found that the base of the jet must be highly magnetized in order not to suffer from devastating radiative energy losses. We also see that large negative values of the reconnection power-law d are effectively prohibited, since reconnection occurring predominantly at the base of the jet will result in heavy radiative losses in the strong magnetic fields. These results are independent of black hole mass. In terms of the equipartition fraction or magnetization, which we define as σ = U B / U e ± , the equation becomes It can be more convenient to use this expression when solving the evolution equations numerically, if U B / U e ± becomes very large. 8.1 Constraining the rate of reconnection in the jet Finally, we try to
Paper: 2021MNRAS.502..351R (fulltext)
left-hand panel of Fig. 3 ). Despite such significant perturbations, the presence of a DM filament potential well coupled to the relatively high density of the gas ensures that the cooling time within the filament remains short. As a result, the filament is still visible as a cold stream cutting through the hot superbubble in the middle right-hand panel of Fig. 3 . Due to these consequent perturbations induced by the stellar feedback, we cannot use either the temperature or vorticity to define the gas filament truncation radii in the feedback run. It should also be noted that our assumption of isothermality of the filaments becomes less valid than in the no stellar feedback case as stronger temperature gradients develop between core and outer envelope. To be more specific, in the no feedback case, the temperature varies between core and truncation radius by about a factor of 2, but in the feedback case in can reach an order of magnitude. However, most of this gradient is localized in the outer parts of the filament, so that the central region retains a significantly large isothermal core. This can be understood by performing the following simple calculation. Neglecting the presence of the wall, we may integrate both the gas and DM density profiles (from equation 4 ) to obtain the filament mass per unit length, μ , and its half-mass radius The fact that the (small) core radius contains half of the mass makes the filamentary material relatively impervious to the stellar feedback/filament interaction: provided the core is shielded from it, there can only be a minor change in the amount of gas mass that the filament carries. It has been suggested in the literature that Kelvin–Helmholtz instabilities could be triggered at the interface between cold filament gas and the feedback powered, hot, galactic wind (e.g. Mandelker et al. 2016 and subsequent work). These will depend non-trivially on redshift and the distance of a filament segment to the central galaxy, so it is quite difficult to define a unique characteristic time-scale, t KH . Nevertheless, writing , where v w is the relative velocity between the wind and the gas filament and ρ w the density of the wind, we can see that given the steepness of the filament density profile that we measure, t KH becomes larger as the perturbation progresses deeper in the filament. This means that the time-scale is ultimately set by t KH ( r 0 ). Plugging in typical numbers for our feedback run at z = 4, i.e. r 0 ∼ 5 kpc, ρ ( r 0 ) ∼ 3 × 10 −27 g cm −3 , v w ∼ 100 km s –1 , and ρ w ∼ 3 × 10 −28 g cm −3 , we, thus, get t KH ∼ 20 Myr which is about an order of magnitude shorter than the infall time from the virial radius of the embedded halo. The conclusion is, thus, that our gas filaments should not survive the interaction. Notwithstanding that this does not happen in our simulations, which might admittedly be of too low a resolution to capture the instability properly, the calculation ignores both the
Paper: 2020ApJ...893...88O (fulltext)
for the jet structure finding almost identical results. Figure 6. GRB 061121 light curves as observed by XRT (energy flux in the 0.5–10 keV band, blue error bars) and UVOT (mean flux density in the white filter, red error bars, de-reddened assuming E ( B − V ) = 0.055) compared to the predictions of our model (solid lines), where both the high-latitude emission (dotted lines) and forward shock (dashed lines) emission are taken into account. The adopted parameters are reported in the text. The w filter model light curve accounts for the transmission curve of the filter, retrieved from the HEASARC website ( https://heasarc.gsfc.nasa.gov/docs/heasarc/caldb/swift/docs/uvot/ ). 3.2.2. GRB 100906A We collected multi-filter optical data from Gorbovskoy et al. ( 2012 ) and A. Melandri (2020, private communication). The observed magnitudes are corrected for galactic extinction (Schlafly Finkbeiner 2011 ) and for extinction in the host galaxy. Figure 7 shows the comparison between the observed light curves and our high-latitude and forward shock emission models, obtained assuming a Gaussian jet structure with Γ c = 160, θ c = 42, θ Γ = 33, E c = 5 × 10 53 erg, n = 15 cm −3 , p = 2.1, e = 0.03, and B = 2 × 10 −2 . For the prompt emission, we assume again a 2BPL spectrum as before, with parameters taken from spectral fitting of the brightest BAT pulse. Taking the fitting parameters from Oganesyan et al. ( 2017 ), we obtain , , and α 1 = −2/3, α 2 = −1/2 as before. The high-energy photon index is set to β = −3.2, and we assume a prompt emission radius R 0 = 3 × 10 15 cm. A second, later emission event with an almost comparable peak flux is present in the XRT light curve at around 100 s post-trigger. The spectrum of this pulse can be fitted by the 2SBPL with E peak ∼ 5 keV and E break ≲ 1 keV, with a high-energy spectral slope β ∼ −3.7. Although it is subdominant, given the softer spectrum, we included the high-latitude emission from this pulse for completeness. For this pulse, we adopt the same jet structure of the first one, and we assume a starting time t 0 = 85 s in the observer frame. Also in this case, the combination of the high-latitude and forward shock emission is able to explain the features of the multiwavelength light curves (see Figure 7 ), with the high-latitude emission dominating over the forward shock emission during the plateau phase in the X-rays, while it is negligible in the optical range. Note that we find equally satisfying results with a power-law structure, setting θ c = θ Γ = 2° and using slopes k E = 2 and k Γ = 3.1. In this latter case, the structure would be very similar to the one we find for GRB 061121, the differences being almost entirely in the distance and ISM density, pointing to the quasi-universal jet structure hypothesis (Salafia et al. 2015 , 2019a , 2019b ). Figure 7. Same as Figure 6 but for GRB 100906A light curves as observed by XRT (energy flux in the 0.5–10 keV band, cyan error bars) and in optical as
Paper: 2016ApJ...825...87R (fulltext)
nearby neighbors or from morphologically different galaxies. If we ignore the contamination (which is significant), the SFR efficiencies determined in Section 3.2.3 would be increased by a factor of about two. This sets an upper limit on the any potential biases in sample selection on the radial surface brightness and the resultant SFR efficiencies. 2.5. Effects of Ly<italic>α</italic> on the Surface Brightness Profiles The Ly α emission line falls into the rest-frame FUV bandpass and could potentially contaminate the measured flux. We consider whether Ly α could be contaminating our stacks, and thereby the resulting surface brightness profiles. Previously, Rafelski et al. ( 2011 ) conducted a test to see whether a typical SFG Ly α line would significantly affect the measured photometry by using the stacked LBG spectrum from Shapley et al. ( 2003 ) to estimate the effect, and found it to be negligible. However, it is possible that a subset of our sample has much stronger Ly α emission than a typical LBG, such as the green peas (Henry et al. 2015 ), which could bias the SFRs obtained from the rest-frame FUV continuum high. To test for this possibility, in particular in the galaxy outskirts, we conduct the following tests. First, for the and samples we create a smaller sample including about half of the galaxies for which Ly α does not fall in the FUV bandpass by constraining the redshift, and compare the radial surface brightness profile of that stack to the full subsample stacks. We find no evidence of a decrease in the flux in the outskirts of the galaxies within their uncertainties. Second, we compare the radial surface brightness profiles for the FUV and NUV bands, since the NUV bands will not include the Ly α line. Again, we find no evidence of a change in the radial surface brightness profile in the outskirts. The fact that the galaxies have the same radial surface brightness profile shapes in both bands suggests that the Ly α line has little to no effect on the radial surface brightness profiles. We note that this is not in conflict with the results by Steidel et al. ( 2011 ), who find extended Ly α halos around SFGs. First, that study samples a different galaxy population of brighter SFGs. Second, the Steidel et al. ( 2011 ) study goes out to ∼80 kpc using ground-based observations, while we limit ourselves to the central ∼10 kpc, which is barely resolved in that ground-based study. Third, Steidel et al. ( 2011 ) use narrowband filters to find the Ly α , while our stacks are in very broad filters, and thus the Ly α would not contribute except in the extreme emitter cases, and then only in the FUV bands for part of the redshift range. 3. SFR EFFICIENCY The radial surface brightness profiles of the SFGs show spatially extended star formation in the galaxy outskirts. Star formation is expected in the outskirts of high-redshift galaxies, as it is measured in atomic-dominated hydrogen gas at low redshift (Thilker et al. 2007 ; Boissier et
Paper: 2021MNRAS.505.3795N (fulltext)
(5) predicted H i mass derived using equation ( 1 ); and (6) calculated H i deficiencies. 3.2 H <sc>i</sc> kinematics of the NGC 7232/3 triplet Although the coarse velocity resolution of our observations does not allow for the detailed analysis of the kinematics of the triplet galaxy group and its surroundings, it is clear from the Fig. 5 that NGC 7232B has higher velocities and is kinematically different from NGC 7232 and NGC 7233 (velocity difference of up to 400 km s −1 ). The H i debris appears to have similar velocities to NGC 7233 and NGC 7232 within ∼200 km s −1 , which could indicate that most of this gas is coming from the interaction between the galaxy pair and not NGC 7232B. Figure 5. MeerKAT-64 H i velocity field map of the NGC 7232/3 triplet. The black stars indicate the optical centres of the triplet galaxies. 3.3 H <sc>i</sc> deficiency equation of the NGC 7232/3 triplet We determine the H i deficiency of each galaxy in the NGC 7232/3 triplet group to determine if the group spirals are H i -deficient or H i -rich. The H i deficiency is derived by comparing the observed H i mass to the expected H i mass for an isolated galaxy. The expected H i masses () are derived from the absolute B -magnitudes obtained from HyperLeda 4 (Makarov et al. 2014 ). We have used the scaling relation derived in (Jones et al. 2018 ) based on measurements or constraints on the H i masses of 844 isolated galaxies compiled by the AMIGA project (Analysis of the interstellar Medium in Isolated GAlaxies; Verdes-Montenegro et al. 2005 ): where log [ L B / L ⊙ ] is derived by where M bol, ⊙ is the Sun's bolometric absolute magnitude, we adopt M bol, ⊙ = 4.88 (Lisenfeld et al. 2011 ), D is the distance in Mpc, and B c is the absolute B -magnitude corrected for extinction. The H i deficiency is then defined by where is the H i mass derived from the observations. Positive values of DEF indicate H i deficiency while negative values indicate H i excess. We compare the expected H i mass for each galaxy to its H i mass derived from our observations (see Table 2 ). NGC 7232 and NGC 7233 are H i -deficient while NGC 7232B is found to have excess H i . Taking into account the uncertainties in B -magnitudes and the scatter in the scaling relation, we consider a galaxy to be H i -deficient when DEF > 0.2 and H i excess if DEF −0.2. The total H i mass of the NGC 7232/3 triplet including all the H i debris (i.e. H i streams and clouds) is ] = 9.84 ± 0.002. Comparing the total derived H i mass to the expected H i mass of , we find a total DEF of −0.23 ± 0.11, which indicates that the NGC 7232/3 galaxy triplet has an overall H i excess compared to field galaxies. 4 DISCUSSION One of the questions coming from our data is the possible origin of the H i debris observed. In this work, we search for signs of interaction within the NGC 7323/3 galaxy triplet and its surrounding environment. The detection of wide-spread H i debris around the galaxy triplet demonstrates the need for