adsabs/WIESP2022-NER huggingface dataset:
Paper: 2018ApJ...854...31J (fulltext)
S = S t (and ). In Table 6 we present the statistical results of a linear least-squares analysis for the toy model ensemble, both prior to and after time-randomization, for , 0.010, and 0.020. These results can be compared directly against the results from Section 4 and Table 4 . As can be seen in the table, for a large value of there is a significant difference in the width of the measured slope histogram between the original time-ordered data and the randomized data, with the distribution of the original time-ordered data being significantly broader than that of the randomized data. In all three model cases, the width of the measured slope for the randomized data remains similar to that for both the time-ordered and randomized data obtained by the JCMT Transient Survey to date, suggesting that the uncertainty assumptions of the toy model are reasonable. As the magnitude of decreases, the measured slope distribution for the time-ordered data becomes narrower. The trends seen here are entirely as expected since the strength of the input slope within the toy model is proportional to and becomes directly measurable in the time-ordered data as the secular signal becomes greater than the uncertainty in the individual measurements. For the randomized data, the input slope is significantly disrupted and thus there is little signal remaining to influence the histogram. The observed trend is similar for the histograms of measured slope in units of the measured slope uncertainty. For large values of this width becomes significantly larger than unity for the time-ordered observations while the randomized distribution remains close to 1.1 (which is also the value obtained by the Transient Survey ensemble). Table 6 Secular Variability Analysis: Histogram Widths in the Toy Model Notes. a Standard deviation of S measured for all members of the ensemble (see the text). b Standard deviation of measured for all members of the ensemble (see the text). The toy model can be extended in order to predict the detection power of the JCMT Transient Survey after its full three years of observation. By increasing the number of observations to 30 and the time period to three years, we show in Table 6 that the difference in ensemble histogram widths between the time-ordered and random observations becomes much stronger for each value of . We thus conclude that at present we can rule out a Gaussian distribution of secular variations linear in time with but that a distribution with would be undetectable today. After the full three years of the survey, we expect to either measure or rule out distributions with . In each of these toy models, the normal distribution of linear slopes prevents there being a significant population of sources with large brightness variations, such as those observed by the JCMT Transient Survey (Table 5 ). This suggests that the five observed sources are indeed rare, statistical outliers. It should be recognized, however, that the toy models are
Paper: 2016ApJ...821...25L (fulltext)
uncertainty for our source, we carried out 2000 ray-trace simulations with MARX 5.1.0 at positions near it and at the same off-axis angle. The spectrum from the multicolor disk (MCD) fit to C1–C3 (Section 3.2 ) was assumed. 2.3. The HST Images There are two images (in two filters: F606W and F814W) taken by the HST /ACS/WFC on 2004 December 15 in the field of our X-ray source. Each image has four frames of 525 s each. As will be shown later in Section 3.1 , these images indicate that our source has a candidate host galaxy SDSS J141711.07+522540.8 (GJ1417+52 hereafter) with extent ∼4′′ and a very faint candidate optical counterpart (sGJ1417+52 hereafter) at a ∼1′′ offset from the nucleus. In order to derive their main photometric parameters, we fitted the HST /ACS/WFC images using GALFIT (Peng et al. 2010 ) with multiple Sérsic components (convolved with the PSF) for the galaxy and a PSF model for the faint optical source, which seems unresolved in these images. In order to improve the statistics, especially for the faint source sGJ1417+5, the fits used 5′′ × 5′′ (centered at the center of the galaxy) stacked images (one for each filter) from aligned and bundled 16 16 http://stsci.edu/hst/wfc3/documents/ISRs/WFC3-2014-24.pdf FLC frames. The effective PSFs 17 17 http://stsci.edu/hst/acs/documents/isrs/isr0601.pdf at the position of our source in the four frames were averaged and used to fit the stacked images. Because of dithering in the observations, different frames had different degrees of distortion at the position of our source, causing some problem in stacking them. However, we found that the fits to single frames with the corresponding PSFs gave similar results. 2.4. The Gemini Spectroscopic Observation The galaxy GJ1417+52 was observed during the night of 2013 February 8 (UT) with the Gemini Multi-Object Spectrograph (GMOS, Hook et al. 2004 ) at the Gemini North Telescope, in queue mode (program ID GN-2013A-Q-37). The data were acquired during dark time (illumination fraction 0.8%), under photometric conditions and ∼070 seeing. The 400 lines/mm ruling density grating (R400) centered at 7000 Å was used. We chose a slit of width 075 and put it through the center of the galaxy GJ1417+52 and the source sGJ1417+52 in order to obtain their spectra simultaneously. A total of four exposures of 1500 s each were obtained. Small offsets in the spectral direction (50 Å) toward the blue and the red were applied between exposures to allow for the gaps between CCDs to avoid any loss of important lines present in the spectra. Spectroscopic flats and comparison lamp (CuAr) spectra were taken after each science exposure. In addition, to derive the sensitivity function and flux calibrate the science spectrum, the spectrophotometric standard star G191B2B was observed as part of the baseline calibration provided by the observation. Because the standard star was observed on a different night (2013 March 03 UT) and under different observing conditions,
Paper: 2020MNRAS.495.1641S (fulltext)
of the X-ray flux (see e.g. Patruno, Wijnands van der Klis 2009b ), which can hinder the detection of relative weak spin derivatives (see e.g. Galloway et al. 2002 ; Burderi et al. 2006 ; Riggio et al. 2008 ; Bult et al. 2019 ). Here, we report on the phase-coherent timing analysis of the AMXP IGR J17591–2342 (Sanna et al. 2018d ), first detected by INTEGRAL on 2018 August 10 (Ducci et al. 2018 ) and extensively monitored by the NICER instrument during its outburst. Moreover, we investigated the evolution of the spin pulse delays within the framework of magnetically disc-threading theories in which the material torque (dependent on mass accretion rate) as well as the disc-threading torques exerted by the magnetic field are taken into account (see e.g. Rappaport et al. 2004 ; Kluźniak Rappaport 2007 ). For the first time in phase-coherent timing analysis of AMXPs, we modelled variations of the frequency spin derivative taking into account instantaneous values of the mass accretion rate by following the X-ray flux evolution. Recently, a similar approach has been successfully applied to the timing of the X-ray pulsar GRO J1744–28, a slowly spinning (∼2 Hz) NS accreting from a low-mass companion (Sanna et al. 2017b ). The obtained results show that, at least in this case, the standard theory of accretion on to NS works very well also for fast rotators and allowed an accurate determination of the magnetic field strength that is in the range of the values expected for this class of sources. Finally, we report on the updated source ephemerides and we discuss the properties of the X-ray pulsation as well as its temporal evolution. 2 OBSERVATIONS AND DATA ANALYSIS NICER (Gendreau, Arzoumanian Okajima 2012 ) started monitoring the X-ray transient IGR J17591–2342 on 2018 August 15 (MJD 58345) up to 2018 October 17 (MJD 58408.2) for a total of 37 dedicated observations. We processed the available NICER observation using the nicerdas pipeline (version V004a) retaining events in the 0.2–12 keV energy range, for which the pointing offset was 54 arcsec, the dark Earth limb angle was >30°, the bright Earth limb angle was >40°, and the ISS location was outside of the South Atlantic Anomaly (SAA). Moreover, we removed short intervals that showed background flaring in high geomagnetic latitude regions. As a result of the filtering process we retained a total exposure time of ∼101 ks. The spectral background was obtained from a blank-sky observation of the RXTE-6 region. We did not detect any Type-I thermonuclear bursts during the observations analysed in this work. For each observation we modelled the energy spectrum with an absorbed blackbody component plus a thermally comptonized continuum ( Tbabs(bbody+nthcomp) in Xspec) and we extracted the source unabsorbed flux in the energy range 0.5–10 keV (see Fig. 1 ). Finally, we barycentred the NICER photon arrival times with the barycorr tool using the JPL DE-405 Solar system ephemeris and adopting the
Paper: 2020ApJ...902L...7Z (fulltext)
field strength can be narrowed by 10 ν Lar max{ ν Rad , ν Em }. Utilizing the fact that Fe ii λ 5362.970 Å () absorption line is fully radiative aligned for all of the orbital phases, we deduce that the upper limit field strength is ∼108 mG. Additionally, we find that the Fe i λ 5586 Å () absorption line presents a more than 3 σ detectability polarization signal at the orbital phases 0.544, 0.564, and 0.988 and that the polarization direction for all those phases aligns with the radiative pumping direction rather than the magnetic field direction. This gives us an upper limit for the field strength of ∼80 mG. The accuracy of the field strength is increased by at least two decades compared to the previous 10 G upper limit, constrained from the non-detection of the Zeeman effect. 5. Discussion GSA is the most probable and natural explanation for the discovered polarization signals of the ground-state absorption lines that are realigned to one direction as opposed to the orbital synchronization. We provide here the mean magnetic field of the photosphere surface medium that corresponds to the central part of the absorption line. Figure 2 (d) illustrates that such magnetic field is reasonable on the stellar surface. A full magnetic field tomography requires further study, which is beyond the scope of current observation precision. Higher-resolution spectral observation will enable us to study the wings in the absorption line profiles, separating the emission components from outflows. Through our analyses, we assume that the anisotropic pumping comes from the secondary (following Leone et al. 2018 ). However, when constraining the magnetic field strength, a narrower range might be achieved if accounting for the photons from the primary. We calculate the extended radiation field of the primary following Zhang et al. ( 2015 ) with the Lambert cosine law and limb-darkening (Claret Bloemen 2011 ) considered, and find that the lower limit for magnetic field strength can be ∼2 mG (see Figure 4 (a)). Additionally, we could interpret the orbital synchronization as a result of local polarized flux illuminating the medium rather than pumping from the secondary. Nonetheless, even under such an assumption, the realigned polarization direction of the two ground-state lines that we report here can still be explained by the magnetic realignment of the angular momentum on the ground state. Moreover, beyond the optical band, GSA can be implemented with multi-frequency data ranging from ultraviolet to submillimeter to trace not only the spatial but also the temporal variations of magnetic field (Yan Lazarian 2007 , 2012 ; Shangguan Yan 2013 ; Zhang Yan 2018 ). 6. Summary Polarizations of absorption lines from the ground state with a high polarization degree are discovered in the binary system of 89 Her. We conclude the following. 1. The polarization directions of the two Fe i lines show little variation across the orbital period as opposed to the polarization of
Paper: 2016ApJ...817..111D (fulltext)
flattening in the values at lower column densities could be an effect of completeness. The nominal function (e.g., blue curve in Figure 9 ) is calculated for 3 σ detections. If this threshold is raised to 4 σ or even 5 σ , there is relatively less effective pathlength at lower column densities. The blue curve in the left-hand panels of Figures 9 , 11 , 12 representing moves to the right by ∼0.1 dex for each 1 σ increase in significance level and the frequency of weak absorbers rises. This results in a steepening of the weak-end slope of the cumulative distribution function (less pathlength for the same number of detections), but it does not eliminate the non-power-law nature of the distributions in any case. Therefore, we believe that the turnovers seen in many of the metal absorber distributions are real, but the weak-end slopes should be treated with caution. 4.2.5. Ne <sc>viii</sc> The extreme-UV (2 s – 2 p ) doublet transition of Ne viii (, f = 0.1030, 0.0505) is often presented as a less-ambiguous tracer of collisionally ionized gas than O vi or other highly ionized, FUV metal transitions (Savage et al. 2005 ). Several compelling studies of warm-hot or multiphase gas have been made using Ne viii in conjunction with H i , O vi , and other ions to constrain gas temperature and density (Savage et al. 2005 , 2011 ; Narayanan et al. 2009 , 2011 ; Tripp et al. 2011 ; Meiring et al. 2013 ; Hussain et al. 2015 ). Unfortunately, the Ne viii doublet only redshifts into the COS/G130M band at , which limits us to the 13 highest-redshift AGN sight lines in this survey, usually without observations of the strong H i counterparts. We measure several dozen Ne viii systems, but all are at redshifts similar to the background AGNs and show the hallmarks of intrinsic, rather than intervening, absorption. They are primarily strong, blended, multi-component absorption profiles in high ions, often lacking H i absorption. There are three Ne viii detections at >3 σ in our survey that can be reasonably identified as intergalactic rather than intrinsic systems. The system detected toward PKS 0405−123 at z = 0.49494 in Ne viii and O vi by Narayanan et al. ( 2011 ) shows and . In an earlier reduction of the data allowing for the contaminating effects of fixed pattern noise, Narayanan et al. ( 2011 ) measured column densities of and . The origin of the Ne viii and O vi is consistent with collisionally ionized gas with T ≈ 5 × 10 5 K and a baryonic column density of cm −2 . Two other Ne viii detections have no corresponding absorption in any other metal or H i transition. The first is a pair of broad absorption features toward PKS 0637−752 consistent with Ne viii doublet absorption at z = 0.60552 with . Ly β and O vi λ 1032 non-detections place column density upper limits of and , respectively. A similar pair of weak absorption features in the HE 0238−1904 sight line is consistent with a z = 0.50511 Ne viii doublet: , , and . These two Ne viii -only systems provide
Paper: 2021MNRAS.501...78P (fulltext)
light. As the equilibrium temperature increases, thermal emission becomes as important as reflected light in that bandpass, leading to a smaller phase curve amplitude since both the dayside and the nightside are emitting light. Then as the equilibrium temperature increases again, the amplitude increases due to the shortening of the radiative time-scale, as in the other phase curves dominated by thermal emission. The third class of behaviour occurs at specific wavelengths that are influenced by exotic chemical effects, where the phase curve amplitude can exhibit non-monotonic behaviour with equilibrium temperature. This occurs in particular in the Spitzer 3.6- and 8-micron bands of Fig. 9 due to methane and CO opacity but will also happen at all wavelengths where a change in chemical composition with temperature leads to a significant opacity variation. For and for , the amplitude of the phase curve rises with equilibrium temperature whereas for the opposite happens. This peculiar behaviour is due to our assumption of local chemical equilibrium that leads to a spatially varying ratio between carbon monoxide and methane molecules on the planet (see also Dobbs-Dixon Cowan 2017 ). Hot Jupiters, however, are not expected to be in local chemical equilibrium since the vertical and horizontal transport of material can be much faster than chemical reactions. In such a case, the system is said to be quenched, and homogeneous chemical abundances are expected at the photosphere (Cooper Showman 2006 ; Visscher Moses 2011 ; Agúndez et al. 2014 ; Drummond et al. 2018a , 2020 ; Steinrueck et al. 2019 ). If quenching happens, and as long as CO is the dominant species (rather than CH 4 ), then the amplitude of the phase curve should be monotonically increasing with equilibrium temperature at all wavelengths dominated by thermal emission. 6.1.2 Phase curve offset The offset of the phase curve maximum is shown in the bottom panels of Fig. 9 . It always follows a bell curve, with an offset that increases with equilibrium temperature at low equilibrium temperatures and an offset that decreases with equilibrium temperature at larger equilibrium temperatures. Whereas the decrease of the offset with increasing temperature is expected based on the radiative and advective time-scales shown in Fig. 4 , the low-temperature behaviour is more puzzling, but is consistent with some of the cases explored in fig. 9 of Komacek et al. ( 2017 ) Again, the Kepler bandpass stands apart. At low temperatures, when the light curve is dominated by reflected light, the amplitude is large and the offset is close to zero because Rayleigh scattering is approximately homogeneous over the dayside of the planet. As the equilibrium temperature increases, thermal emission becomes more and more important compared to Rayleigh scattering, leading to an offset that increases with increasing equilibrium temperature. Above , the decrease of the radiative time-scale with equilibrium temperature
Paper: 2020A&A...639A..62K (fulltext)
left. Fig. 2 Schematic overview of the features in the outer circumbinary ring detected in our SPHERE PI image. The image is centered on the location of GG Tau Aa and was smoothed for illustration purposes. Any unresolved circumstellar material may create a halo of polarization signal around the star. Because this unresolved polarized signal can affect the analysis of the immediate circumstellar environments, we subtracted these polarized signals (i.e., the total intensity halo multiplied by the degree of polarization) individually for Aa and Ab. While subtracting the polarization signal of Aa slightly increases the contrast of the fine structures in the immediate stellar environment, subtracting the polarization signal of Ab instead blurs these structures. This can be explained by the fact that the measured polarization degree of Ab is somewhat higher than that of Aa. Therefore, subtracting the polarization signal of Abadds an artificial polarization halo around Aa, which weakens the fine structures in its environment. Subtracting the (less strongly) polarized signal of Aa, however, does not noticeably affect the environment of Ab. Figure 3 shows the resulting image after subtraction of the polarization signal of Aa with annotations of the detected features. The immediate stellar environments that are affected by the diffraction pattern are masked out. In all the images, the inner region appears highly structured, as highlighted in Fig. 3 by the dotted lines. Most prominently, the “north arc”, an extended structure to the northeast of Ab observed in previous scattered light images (e.g., Krist et al. 2002 , 2005 ; Itoh et al. 2014 ; Yang et al. 2017 ), is clearly detected and appears in our SPHERE image to be composed of a double-arc structure at projected separations of ~ 0.38′′ and ~0.48′′. This double-arc structure may extend along the entire eastern side to the south, interrupted by a dark lane extending from Ab toward the east (see the dotted gray lines in Fig. 3 ). This dark lane seems to be connected to the shadow observed in the outer disk at a similar position angle (see Sect. 4.4 ). We furthermore detect two additional filament structures northwest and southwest of Ab. It is unclear, however, whether they are related to the double-arc system on the eastern side. Finally, another filament is detected immediately southeast of Aa, pointing toward the south. Fig. 3 Zoom on the inner region after subtraction of the stellar polarization on Aa. The dotted lines highlight the detected features: the “north arc”, revealing a double-arc structure (green), several filaments (light blue) and a possible shadow lane (gray dashed). See Sect. 4.1 for details. The immediate stellar environments (120 mas) are masked out. North is up and east is to the left. Figure 4 (left panel) shows the angles of linear polarization overplotted on the inner disk region. The polarization angles θ were calculated according to θ = 0.5 ×arctan( U ∕ Q ), within bins of 3
Paper: 2016ApJ...817..106B (fulltext)
smaller than the weaker CO lines in the planet transmission spectrum. Furthermore, in Doppler space the nature of this spurious signal is substantially different from the planet absorption, and largely canceled prior to cross-correlation. Winn et al. ( 2007 ) demonstrated that HD 189733 is a slow rotator with a rotational period close to 13 days. Thus any configuration of spots can be considered static during one single transit observation (less than 2 hr). Compared to the observer, the star moves in radial velocity by only 0.2 km s −1 during the transit (mostly due to the change in barycentric velocity), which is only 1/15 of the full width half maximum of the CRIRES IP. This means that any distortion of the stellar lines caused by unocculted spots is static in both radial velocity and time. In this case our data analysis designed to remove telluric lines (Section 3.4 ) will also re-normalize residuals from uncorrected star spots, further dampening the already small signal. However, we know from past transit observations (Pont et al. 2007 , 2008 ) that HD 189733 b usually crosses some spots during transit. In this case, the distortion of the stellar line suddenly disappears (when the spot is occulted) and reappears, generating a pulse radial velocity signal in the data, uncorrected by our analysis. The duration of this pulse is dependent on the size of the crossed spot (or region of spots). Although probably unrealistic, we discuss the worst-case scenario in which all spots are clustered in a region of radius 0.1 R S (based on the 1% coverage and spots significantly darker than the photosphere). For an impact parameter of b = 0.66, the length of the chord crossed by the planet is 1.5 stellar radii. With 39 frames observed in transit, the planet crosses 0.038 stellar radii per frame. This unrealistically big group of spots would therefore affect 2.6 frames of our data. For a spot signal 1/10 of the planet signal, occulted star spots can only contaminate 8% of the detected CO signal, or based on the detection presented in Section 5.1 . We therefore conclude that—at least in the case of HD 189733 b—the contamination due to uncorrected star spots is negligible. 3.4. Removal of Telluric Lines Telluric lines change in depth during the observations, but not in position. Therefore, the telluric spectrum is most effectively removed by modeling the flux of each data column (i.e., each wavelength channel) in time. The main cause of telluric variation is the change in geometric airmass. Changes in the IP (due to, e.g., in-slit guiding, seeing, and performances of the adaptive optics), as well as changes in the water vapor content in the Earth’s atmosphere could also affect the flux in the core of telluric lines. We initially normalize each spectrum by its median continuum level. We then model the time-variation of the flux in each spectral channel via the simple functional form:where A is the geometric airmass, and we divide through the fit. The