adsabs/WIESP2022-NER huggingface dataset:
Paper: 2020MNRAS.494.3699X (fulltext)
An additional effect, therefore, will be required to explain the orbital variation of the GeV emissions of LMC LP3 and especially for 4FGL J1405.1−6119. In addition to the effect of the anisotropic soft-photon field, we will explore the orbital modulation caused by the dependence of travel distance of the unshocked pulsar wind that moves toward the observer on the orbital phase. This effect is introduced in Sierpowska-Bartosik Bednarek ( 2008 ) for the PSR B1259−63/SS2883 system and in Sierpowska-Bartosik Torres ( 2008 ) and Khangulyan, Aharonian Bosch-Ramon ( 2008 ) for the very high-energy emission (TeV emission) of LS 5039. Since the interaction between the pulsar wind and stellar wind/disc will produce a termination-shock with cone-like structure (Canto, Raga Wilkin 1996 ), the distance from the pulsar to the termination shock ( r s in Fig. 2 ) increases with the angle measured from the axis created by the pulsar and the companion star. This causes a dependence of the distance from the pulsar to the termination shock in the direction of the observer on the orbital phase and results in the orbital dependent intensity of the ICS of the unshocked pulsar wind. For the gamma-ray binary with O-type star, the GeV light curve is not well studied with the effect of the shock geometry and the system parameters. In this paper, therefore, we will perform an investigation for the modulation of the GeV emission with the shock-cone model, and will apply the model to two gamma-ray binaries, LMC P3 and 4FGL J1405.1−6119. LMC P3 is the first gamma-ray binary outside Milky Way and it is discovered in the Large Magellanic Cloud (LMC) (Corbet et al. 2016 ), which is a neighbour galaxy located at about 50 kpc from the Earth (Macri et al. 2006 ; Pietrzyński et al. 2013 ; de Grijs, Wicker Bono 2014 ). The companion star in the LMC P3 system is an O-type star, but the compact object has not been identified yet. In the LMC, ∼50 supernova remnants (SNRs) are known to emit X-rays (Seward et al. 2012 ), and LMC P3 is associated with SNR 0535−67.5 in the H ii region DEM L241 (Davies, Elliott Meaburn 1976 ). Long, Helfand Grabelsky ( 1981 ) report the first detection of the X-ray emission from the SNR 0535−67.5. Seward et al. ( 2012 ) find a compact X-ray source (CXOU J053600.0−673507) with a luminosity of ∼2–3 × 10 35 erg s −1 in 0.3–10 keV energy bands and an optical counterpart, O5III(f) star. They therefore suggest that CXOU J053600.0−673507 is the high-mass X-ray binary hosting neutron star or black hole. Corbet et al. ( 2016 ) analyse the GeV counter part (LMC P3) of CXOU J053600.0−673507 in data of Fermi Large Area Telescope ( Fermi -LAT) and find a 10.3-d periodic modulation, which reveals that LMC P3/ CXOU J053600.0−673507 is the binary system with a orbital period of P orb ∼ 10.3 d. They also confirm the orbital modulation of the X-ray emission. HESS Collaboration et al. ( 2018 ) subsequently detect the TeV emission from LMC P3. van Soelen et al. ( 2019 ) find
Paper: 2018ApJ...869...44K (fulltext)
is more interesting. Even when the initial photon energy lies well below the first Landau resonance, the scattering cross section spikes (but does not diverge) at a value of ω f (and therefore ω i ) somewhat larger than 2 m . This spike arises from the pole in the matrix element associated with annihilation of the initial electron with a virtual positron into the final-state photon (Herold 1979 ; Daugherty Harding 1986 ). It appears at a lower photon frequency than the first Landau resonance (energy ) if B > 4 B Q . Figure 3. Differential nonresonant electron scattering cross section vs. initial photon frequency ω i (top panels) and final photon direction cosine μ f (bottom panels), in the regime B ≫ B Q . Curves in the top panels represent . Black curves: μ f > 0; blue curve: μ f = 0; red curves: μ f 0. The cross section rises as the scattered photon approaches the threshold for single-photon pair creation. In the bottom panels, the blue curve marks the classical limit, and black and red curves are used for clarity. Dotted curve: the final-state photon rapidly converts to an electron–positron pair. The partial cross section for scattering into non-pair-creating states is shown in Figure 4 . The high-energy behavior of the total cross section is opposite to the Klein–Nishina result for scattering in vacuum: the cross section grows as ω rises above m . Nonetheless, scattering is still suppressed for photons that have a large energy owing to relativistic motion of the scattering charge along B (Lorentz factor γ ≫ 1), because the factor decreases as ∼1/ γ 2 . Figure 4. Integral nonresonant electron scattering cross section plotted vs. initial photon frequency ω i for a range of initial direction cosine μ i , in the regime B ≫ B Q . The high-frequency peak marks the opening up of pair conversion in the final state (see Figure 6 ). The high-frequency dip in the curves shown in Figure 4 represents the opening up of the final-state phase space to pair creation over some range of scattering angles. A scattered photon can convert directly to an electron–positron pair if , meaning that the cross section for scattering-assisted pair creation, e ± + γ → e ± + e + + e − , can substantially exceed the vacuum value. This phenomenon is examined in more detail in Section 4 . More generally, a high-frequency Klein–Nishina scaling for the scattering cross section—which is approached in the case of scattering at high-order Landau resonances in sub-QED magnetic fields (Gonthier et al. 2000 )—loses meaning as B rises above 4 B Q , because the scattered photon has a high probability of converting to a pair (Beloborodov Thompson 2007 ). 3.2. Finite-<italic>B</italic> Correction to the Cross Section A simple correction to the electron–photon scattering cross section representing a finite magnetic field is available. The overlap of the photon wave function (wavevector k μ ) with a pair of harmonic oscillator wave functions, such as appears in the scattering matrix,
Paper: 2018ApJ...869..119E (fulltext)
function at the low end. If the space density of metal-poor GCs is ∼2 Mpc −3 , as estimated by Boylan-Kolchin ( 2017 ) using the correlation between GCs and dark matter halos, then the associated stellar density is about half the above value, considering also the metal-rich GCs that formed later. Evidently, a high fraction of star-forming regions included at least one GC before about half the age of the universe. A similar calculation illustrates the importance of GC formation to the universal SFR at early times. If we consider that the duration of the star formation event that made a GC was Δ t , and the redshift range over which they formed was from z = 7 to 2, a period of 2.6 Gyr, then the average SFR from GC-forming events with M max ∼ 10 6 M ⊙ is (5 × 10 7 M ⊙ /Δ t ) × (Δ t /2.6 Gyr) × (8 Mpc −3 ), where the second term is the fraction of time each event occupies in the GC-forming era. The resultant average associated SFR is 0.15 M ⊙ yr −1 Mpc −3 , which is the same as the peak value of the comoving cosmic SFR density in the universe that occurs at redshift ∼2 (Madau Dickinson 2014 ). These two calculations suggest that the early universe was filled with star-forming regions containing young massive clusters (YMCs) that ended up as today’s GCs. There was little additional star formation making smaller clusters or OB associations that dispersed, aside from the smaller clusters and associations directly connected with the GCs. Today the mass density of GC stars is their space density times their average mass of ∼2 × 10 5 M ⊙ (Harris 2001 ), or 1.6 × 10 6 M ⊙ Mpc −3 . This is 0.26% of the current total stellar mass density from Madau Dickinson ( 2014 ), a factor of ∼1/380. The first factor of 50 in this 380 is presumably from scattered stars that formed with the GCs, and then the remaining factor of ∼8 would be from stars that formed recently without making massive clusters. We can also learn about early star formation in dwarf galaxies from their old stellar populations. WLM has an old metal-poor GC whose luminosity suggests that it contains ∼10 6 M ⊙ of stars (Elmegreen et al. 2012 ). The mass of other stars at the same age and metallicity is comparable to this (Leaman et al. 2012 ; Larsen et al. 2014 ), suggesting that the GC was a high fraction of the total WLM mass when it formed (Elmegreen et al. 2012 ). Similarly, the GC mass in the Fornax dwarf spheroidal was ∼0.25 of the galaxy stellar mass when the GC formed (Larsen et al. 2012 ). For IKN and NGC 147, these fractions are ∼0.5 and 0.06, respectively (Larsen et al. 2014 , 2018 ). These are enormously high fractions, especially considering the additional cluster mass that might have been present in the full cluster mass function. Perhaps the clustered fraction of stars at birth was larger than the 25% assumed above, even close to 100% as suggested by Messa et al. ( 2018 ) for regions with high SFR densities. In either case, the star formation events that formed the observed
Paper: 2021MNRAS.500..520R (fulltext)
and smaller p -values. Figure 4. The slope of the power-law fit to the d N /d m p mass distributions over time. The power-law index p is calculated according to equation ( 16 ). Top panel: the data for the L02 simulations. Bottom panel: the data for all seven simulations. As in Fig. 3 , when computing the power-law index, every data point in all of the L02 simulations is included in the calculation for p , and the same goes for the L04 data. Both panels: The shaded region represents σ , the error in p , calculated according to equation ( 17 ). In the bottom panel, when the larger domains are considered and the planetesimal populations from the multiple L02 and L04 runs are combined, there is much less variance in the value for p . The steady, decreasing trend with time is readily apparent. At t = 50Ω −1 , when enough planetesimals have formed to compute a reliable value for p , the values range between 1.6 and 2.0 across the different sized simulations, and well after planetesimals have formed, at t = 150Ω −1 , the values are between 1.4 and 1.6. A decrease in p represents a shift towards fewer and more massive planetesimals at late times. A trend toward larger masses with time is somewhat expected. However, a clear demonstration of this trend has not been demonstrated in previous studies. In fig. 3 from Simon et al. ( 2017 ), the authors show data for p over time in their simulations, but only over a relatively narrow range of time. 4 Similarly, Schäfer et al. ( 2017 ) show how the values of their fit parameters change over time, but also only for a narrow window. This decrease in distribution fit parameters emphasizes that care is needed when attempting to extract a single value for the power-law index or a single set of parameters that describes the mass distribution of planetesimals formed by the SI. The mass distribution is transient and should be expected to evolve indefinitely, albeit as a slowing rate, particularly when a larger simulation domain provides for more material as shown in the next section. 3.3 Total mass of dust in planetesimals and the onset of planetesimal formation Fig. 5 shows the total mass of the planetesimals in the simulations over time. As in Fig. 2 , the larger domain simulations are divided into smaller subdomains with the same area as the L02 runs. The L08 run shows the largest variance in these data. After t = 40Ω −1 , the spread in the total dust mass in planetesimals in any of the subdomains from the L08 run spans 5– of the total mass of dust in a single subdomain. The data from the L02 and L04 runs generally fit within the maximum-minimum bounds of the L08 run. Once again we note that simply re-running these simulations with a different random seed leads to significantly different consequences for planetesimals formation, shown here directly by the wide spread denoted by the red shaded region that is quite similar to the region-to-region variation in the larger domains. Figure 5. The total mass of the
Paper: 2019MNRAS.484.5578R (fulltext)
4026, Si iv 4089, He i 4144, He ii 4200, He i 4387, He i 4471, He ii 4542, Si iii 4552, He ii 4686 and He i 4713. From the observed ratio values, we used the tabulated and graphical results of Martins ( 2018 ) to define spectral types and luminosity classes, as well as, if necessary, the associated qualifiers. When most of the measured ratios indicate the same result, the associated spectral type is then assigned to the star. However, when a satisfactory agreement is not obtained, the distinct types are then used as initial estimates, and a better tuning on the spectral type of the star is then obtained from the comparison of the observed spectrum with those of standards. Figure 1. Normalized blue-optical spectra (4000–4950 Å) of the newly identified O-stars in the LAMOST DR3. The main spectral features used in the blue-optical canonical spectral classification process are labelled. LS V +39 27 (O4 V), together with J20261845+4347472 (O5 III) and TYC 3339-851-1 (O5 V) are the earliest exemplars in our sample, which is complemented by several mid- to late-type O-stars. Figure 2. Normalized blue-optical spectra (4000–4950 Å) of the newly identified O-stars in the LAMOST DR3. The main spectral features used in the blue-optical canonical spectral classification process are labelled. LS V +39 27 (O4 V), together with J20261845+4347472 (O4 IV(f)) and TYC 3339-851-1 (O5 v ((f))) are the earliest exemplars in our sample, which is complemented by several mid- to late-type O-stars. 3.1 Spectral types and luminosity classes In order to obtain the EWs and full width at half-maximum (FWHM) of the spectral lines necessary for the spectral typing process, we followed standard procedures briefly described here. The line-profile parameters were measured using several spectroscopic analysis routines available in the iraf 2 splot package. For each spectral line, we performed model fittings using both Voigt and pure Gaussian profiles, with the associated continuum computed from regions far from the observed line wings (in general, several tens of angstroms from the estimated line centres), with the final uncertainty in the EW and FWHM values ranging from about 10 per cent for spectra with the highest S/N (usually above 300–350) to 15–20 per cent for the spectra presenting lower S/N values. Table 2 presents the EW values measured from all spectra of stars in Table 1 , for the following absorption lines: (i) eight He lines corresponding to the He i-ii λ4026, He i λ4121, He ii λ4200, He i λ4387, He i λ4471, He ii λ4542, He ii λ4686 and He ii λ4713 transitions; (ii) two Si lines corresponding (when present) to the Si iv λ4089 and Si iii λ4552 transitions. Table 2. Equivalent widths (Å) and associated errors of selected transition lines measured in the LAMOST DR3 O-star spectra. From inspection of Figs 1 and 2 , and from the measured parameters for the He ii lines, we can see that the earliest O-stars in our sample are stars 10, 24, 3, 8, 14 and 19. By following the
Paper: 2017ApJ...848..131Q (fulltext)
blue dots represent the brightness change during the outburst, while the green dots refer to the brightness outside of the outburst. 3. Discussions and Conclusions During the outburst of J162117, the spectral observations showed some properties of CVs, i.e., the broad, two-peaked emission lines (FWHM corresponding to 1500 km s −1 ), , and He ii 4686. These are indications of the accretion on a white dwarf (Scaringi et al. 2016 ). Moreover, the strong rotational disturbance of the emission lines in the eclipse indicates the presence of a rapidly rotating disk (Thorstensen 2016 ). Therefore, after ruling of the merging of a contact binary as an explanation of the outburst of J162117, this event was attributed to a dwarf nova outburst (Scaringi et al. 2016 ; Kjurkchieva et al. 2017 ). However, in comparison with known CVs, J162117 shows several peculiarities (Kjurkchieva et al. 2017 ), namely a deeper eclipse at the outburst than at the quiescence state (by a factor 2.8) and outburst amplitudes at the lowest limit of dwarf nova eruptions. These properties reveal that J162117 is not a normal CV. The binary system was monitored photometrically from 2016 April 19 to September 14, using several small telescopes in the world. We were lucky to obtain several high-precision light curves before the outburst. As shown in Figure 2 , there are two shoulders around the primary minima in the B - and V -band light curves. The two minima show a clear U-type shape and that their bottoms are nearly flat. The color curves plotted in the figure also show a U-type shape around the primary minima, while the color curves are flat outside the eclipses. All of these properties suggest that there are no accreting disks around the white dwarf and thus, the white dwarf is not accreting material from the cool secondary at the pre-outburst quiescent state. Apart from the eclipse of the white dwarf component, the BVR light curves shown in Figure 2 are dominated by ellipsoidal variability in a binary star as well as the magnetic activity of the K-type secondary. As displayed in Figure 2 , some small flare-like events are visible in the B -band light curves, while these events are not seen in the V and R band light curves. These are general properties of optical flares (e.g., Qian et al. 2012 ). Moreover, as shown in Table 1 , the errors of the B -band magnitudes for J162117 and the comparison star are about ±0.006, indicating that those B -band flares are true. The light curves in the figure are asymmetrical, displaying the O’Connell effect. Meanwhile, the depth of the secondary minimum (at phase 0.5) is deeper than that of the primary one (at phase 0.0). These observed properties could be explained by the secondary being covered by dark spots near the inner Lagrange point L 1 where the mass is expected to flow onto the primary. The fact that regions near the L 1 point are heavily spotted has been observed in the pre-CV V471 Tau and in a few CVs (e.g., Hussain et al. 2006 ;
Paper: 2019MNRAS.490.4985S (fulltext)
3 arcmin, and a zoom-in of a 45 arcsec region for closer inspection. The generated image is reasonably convincing as an extragalactic deep field from the point of view of cursory visual inspection, although like many GAN-generated images, on closer inspection there are artefacts (such as low surface brightness features, pixelization, and discontinuities not seen in the real image) that betray the counterfeit. These artefacts are likely caused by a lack of resolution in the GAN generator, and could be remedied through the addition of more layers, or through the use of a larger convolutional kernel. Both of these approaches would increase the receptive field. The GAN also cannot reproduce some of the larger galaxies with the same level of detail as the real image. Nevertheless, the generated image does contains a mixture of early- and late-type galaxies amongst a more numerous field of background sources. The early and late types have the correct colour and morphology, i.e. classic red/elliptical and blue/spiral, respectively, and background sources have the clumpy and disturbed morphology typically seen in the high-redshift galaxy population. Fig. 3 presents single-band ( F 775 W ) thumbnail images of a sample of 100 randomly chosen sources with F 775 W AB magnitudes in the range 21–26 mag. 50 sources are selected from the real image and 50 are selected from the generated images, but the thumbnails have been randomly shuffled in the figure to demonstrate to the reader the visual similarity of individual galaxies in the generated and real data. Figure 2. Example of a generated deep field. The top left panel shows a generated deep field with the RGB channels corresponding to the F 814 W , F 775 W , and F 606 W bands. The size of the generated image was set to 3.2 arcmin, spanning the full XDF (top right), which has a 60 mas pixel −1 scale. The lower panels show 45 arcsec zoom-ins of the generated and real fields. Figure 3. Thumbnail images of real and generated ‘galaxies’. Each image is 7.2 arcsec across and shows the F 775 W channel on an identical linear grey-scale for 100 sources selected in the range 21–26 mag. 50 targets are selected from each of the real and generated catalogues and displayed in random order in row-wise ascending order of magnitude. To test whether the similarity between the generated and real image is more than superficial, we extract ‘galaxies’ from the generated images and compare to those in the real XDF. We use the source finder sextractor (Bertin Arnouts 1996 ) to detect galaxies in the combined F 606 W + F 775 W + F 814 W (pixelwise-mean) images, measuring the corresponding source flux in the individual bands. Sources are identified with the criteria that five contiguous pixels have a signal ≥5σ above the local background. The same detection procedure is applied to the real and generated images, and the photometric zero-points for the generated images are identical to the real data. Fig. 4 compares the magnitude
Paper: 2018MNRAS.477.1356Y (fulltext)
features and high brightness temperatures of the compact radio core component on pc scales (e.g. Lal, Shastri Gabuzda 2004 ; Giroletti Panessa 2009 ; Doi et al. 2013 , 2015 ). Generally, the X-ray emission in radio-quiet AGNs is believed to be produced in the hot corona by the inverse Compton scattering of the soft photons from the accretion disc (e.g. Haardt Maraschi 1991 , 1993 ; Liu et al. 2002 ), while the origin of the radio emission in the radio-quiet AGNs is still not fully understood, which is suggested to be from jets by some works although not fully proved (e.g. Czerny You 2016 ). It is shown that the study of the relation between the radio luminosity and X-ray luminosity is a very useful tool to investigate the coupling of the accretion and the jet (Merloni, Heinz di Matteo 2003 ). The radio/X-ray correlation has been investigated for different types of AGN, as well as black hole X-ray binaries (BHBs) in the low/hard spectral state (e.g. Merloni et al. 2003 ; Wang et al. 2006 ; Li et al. 2008 ; de Gasperin et al. 2011 ). It is found that in the low-luminosity AGNs (LLAGNs) and in the underluminous low/hard spectral state of BHBs (e.g. Körding, Falcke Corbel 2006 ), while this correlation becomes significantly steeper, i.e. , in the luminous AGNs and the luminous low/hard spectral state of BHBs (e.g. Corbel et al. 2004 ; Rushton et al. 2010 ; Dong, Wu Cao 2014 ). Some NLS1 have been included in previous studies of radio/X-ray correlation. However, as the number of the NLS1 in the sample is very small, their statistical properties are overwhelmed by other types of AGNs (e.g. Merloni et al. 2003 ; Dong et al. 2014 ; Panessa et al. 2015 ). Furthermore, the measurements of their radio luminosities could be severely contaminated by the host galaxies; e.g. if the resolution of the radio observation is not high enough, the star formation activities in their circumnuclear regions can have very strong contributions to the radio luminosity (e.g. Deo, Crenshaw Kraemer 2006 ; Ohta et al. 2007 ; Sani et al. 2010 ). Caccianiga et al. ( 2015 ) found that the star formation activities can contribute a significant fraction of radio emission even in some radio-loud NLS1. In order to reduce the contaminations from the host galaxy as much as possible, in this work we make use of radio fluxes within the central regions () of nearby radio-quiet NLS1 measured from high-resolution radio images to study their nuclear radio/X-ray relation, which was rarely studied before. This could help us understand the physics of the accretion and jet in high Eddington ratio regime. We also explore the mutual relation between the radio luminosity, X-ray luminosity, optical luminosity, and black hole mass. There are several works on the radio images of radio-loud NLS1 on kpc scales or even pc scales recently (e.g. Doi, Asada Nagai 2011 ; Doi et al. 2012 ; Gu et al. 2015 ). But we do not choose the radio-loud ones so as to avoid the relativistically beaming effects.