adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016MNRAS.458.2303S (fulltext)
that (except for initially small pitch angle cases) sin θ o ≫ 1/γ o or γ ∥ /γ o ≪ 1 and, from equation ( 17 ), we can write This implies that for τ = γ ∥ /η ≈ 1/(ηsin θ o ) we have γ ≈ 1.3/sin θ o . Thus, even if an electron had started with an almost infinite energy, it loses most of its kinetic energy in a time interval of the order of 1/η, reducing to perhaps a mildly relativistic status (for not too small an initial pitch angle). For instance, let us consider γ o = 10 3 and θ o = π/4; then γ ∥ ≈ 1/sin θ o = √2. Then from equation ( 17 ) or equation ( 19 ) we obtain, for τ = 1/η, γ = 2.3. In another example, taking γ o = 10 4 and θ o = π/3, for τ = 1/η we obtain γ ∥ ≈ 1/sin θ o = 2/√3 and γ = 1.7. Thus 1/η represents the characteristic decay time of synchrotron electrons over which they turn from ultrarelativistic into mildly relativistic ones. On a shorter time-scale, where γ is still highly relativistic, we can write Then we have Equation ( 21 ) is the same relation as in equation ( 3 ), which tells us that there is an upper energy cut-off for highly relativistic electrons in a synchrotron source, . Also, while τ 1/2 = 1/(ηsin 2 θ o γ o ) gives the half-lifetime of the synchrotron electrons (Kardashev 1962 ; Pacholczyk 1970 ; Rybicki Lightman 1979 ), τ ∼ 1/η gives the decay time over which radiating synchrotron electrons turn from ultrarelativistic into mildly relativistic ones. This follows from equation ( 17 ), which gives an exact formula true for all γ, including mildly relativistic electrons, and includes the effects of the pitch-angle changes implicitly. An explicit expression for the pitch-angle changes is derived in the next section. 4 PITCH ANGLE CHANGES Both β and θ in β ⊥ = β sin θ are functions of time. Therefore we can rewrite equation ( 13 ) as Also, since , we have Eliminating from equations ( 22 ) and ( 23 ), we obtain (also see Petrosian 1985 ) This is the relation for the rate of change of the pitch angle of a charge undergoing synchrotron radiative losses. The negative sign implies that the pitch angle decreases with time and aligns with the magnetic field. The rate of alignment is very slow for low pitch angles (θ ≈ 0) as well as for high pitch angles (θ ≈ π/2) and the highest rate of change of pitch angle is for θ = π/4. With the help of equation ( 16 ), we can integrate equation ( 24 ) to obtain In equation ( 25 ), θ θ o , because the pitch angle always reduces with time. There are many notable points. If θ o = π/2, then θ = π/2 also, because if the pitch angle is π/2 then the radiating electron always moves in a circular path in the plane perpendicular to the magnetic field. Also, if θ o = 0 then θ = 0 too, as there is no more reduction in pitch angle. For any 0 θ o π/2, θ → 0 as τ → ∞. For large γ o values, which can be used to estimate the change in pitch angle with time. For example, for, say, θ o = π/3 and θ = π/6, cosh (ητsin θ o ) = 3, which gives τ ≈ 2/η for this change in the pitch angle. Thus there are
Paper: 2017ApJ...840...54C (fulltext)
region creates a bright patch between n0 and p1, and hot plasma escapes out along the new open field to become the spire of Phase 1. The spire is collimated by the dominant strong open magnetic flux. Heat from the reconnection is magnetically confined to a small volume in the closed field below the reconnection location and so shows up as a brightening, but in the open field above the reconnection site it is spread over a much larger volume and so does not lead to a brightening in region n1. The continuous magnetic cancellation makes the jet move sideways. It can be seen, by comparing Figures 7 (b) and (c), that the location of the open field line that has just reconnected (and therefore the spire) moves to the left as reconnection proceeds. This transverse motion can be faster than the photospheric cancellation speed, since the transverse motion of photospheric sources can lead to a buildup of a current sheet that subsequently reconnects. Phase 2 is driven by the subsequent cancellation between n0 and p2 (Figure 7 (d)) and produces a spire and a brightening in the same manner as in Phase 1. It is possible that the two phases are independent of one another, but this is unlikely since they occur so close in time and space and the spire has a more or less smooth motion, which suggests that the first phase helped trigger the second phase or that the fragments p1 and p2 are not independent but are linked somehow by a common convection pattern in the moat below the photosphere. Figure 7. Schematic of the topology and magnetic reconnection in the two successive phases of a complex jet showing a few representative field lines marked by light-headed arrows before and after bursts of reconnections. (a) Before reconnection n0 and the outer parts of n1 and n2 are open magnetic fields with negative polarity. On both sides of n0, there are closed magnetic fields with positive polarities p1 and p2 linking to the inner parts of the negative polarities n1 and n2, as shown. Arrows mark the direction of magnetic field lines. (b) Phase 1 of the reconnection process driven by magnetic cancellation between n0 and p1. Reconnection takes place at the X-point, with plasma velocities shown by solid-headed arrows. The upward-moving outflow jet from the reconnection site is channeled upward by the open field lines as a spire. The separatrices and the field lines that have reconnected (a new open field line and a new closed field line) are marked in red. The location of a chromospheric brightening is also indicated below the reconnection location and straddling n0 and p1. (c) Phase 1 continues with more reconnection at the X-point and so creates a lateral motion of the spire. (d) Phase 2 repeats the same process, but at the X-point on the left side of n0, driven by magnetic cancellation as p2 approaches n0. The separatrices and the newly reconnected field lines are shown in red. The formation and evolution process of the jet is a process in which magnetic
Paper: 2018MNRAS.473.3700M (fulltext)
2015 ). For illustration, in Fig. 3 , we plot the radial profiles of various quantities for two models: inside-out scenario ( Case 1) and for constant short time-scale gas infall rate t f = 2.25 Gyr using the boundary τ b = 100 Myr. As we see, the distributions do not differ significantly for the two shown examples. Notice that, although near the corotation the reduced SFR for high-mass stars , their sum does not tend to zero at the corotation. 3 Our computed profiles for μ g , μ s and are close to observed ones derived by other authors (cf. our Fig. 3 with fig. 1 in Yin et al. 2009 or fig. 4 in Kubryk et al. 2015 ). Figure 3. Radial distributions of various quantities at present time. Solid lines: model with constant t f = 2.25 Gyr. Dashed lines: inside-out scenario. The boundary τ b = 100 Myr. The bend in μ g ( r ) at 4 kpc, seen for inside-out scenario (dashed lines), is due to change in t f approximation for r ≤ 4 kpc. The feature in at this distance is due to (see Sections 3.1 and 3.2). The upward jump in μ g , at r ≈ 12 kpc for the model with constant t f (solid lines) is associated with the outer Lindblad resonance r L , which is located at the same distance (see Section 3.3). The jump maybe easily explained by the cut-off at the same distance. 3.4 Oxygen synthesis in the Galactic disc In general, our ‘chemical’ equation, which describes the radial evolution of oxygen abundance along the Galactic disc in ISG, is close to the classical ones proposed by Tinsley ( 1980 ), Lacey Fall ( 1985 ), WS , PC , Matteucci ( 2003 ) and others with one correction that we work with composite SFR (see equation 8 ): where μ O is the mass density of oxygen in the disc, Z O = μ O /μ g , is the enrichment rate of ISG by newly synthesized oxygen: correspondingly P O is the mean mass of newly synthesized oxygen ejected per CC SN event. The lower limit in the integral, entering equation ( 18 ), we adopt to be equal to 10 M ⊙ since exploding CC SNe do not produce oxygen if the initial masses of their progenitors are less than 10 M ⊙ (e.g. Wanajo et al. 2009 ). The meaning of various terms in equation ( 17 ) is as follows. The second term on the right-hand side of the equation describes the increase of oxygen density due to intergalactic gas infall on to the disc. The third term describes the rate of oxygen return into ISG due to wind from low-mass stars and their pass to remnant; the fourth, fifth and sixth terms describe the depletion of oxygen in ISG due to capture it by the forming objects. The last term in equation ( 17 ) is the radial diffusion of oxygen that we associate with turbulence in ISG. The inclusion effect of diffusion was provoked by the following: if we do not take into account the diffusion, in our model we would have deep gap in oxygen radial distribution near the corotation because here ψ H → 0. To smooth away this sharp gap in the distribution, we must include the effect of diffusion. The combined effect of corotation resonance and
Paper: 2020ApJ...901...19G (fulltext)
inside the of each cluster. Many clusters have a so scarcely populated RGB that this parameter turns out to be null, which essentially points to negligible crowding in the cluster core compared to the environment of typical field-AGB stars. More populous clusters have a significant number of RGB stars, but their never exceeds 1.1 . At these maximum RGB densities, the magnitude bias turns out to be tiny (about mag), although the magnitude rms error (bottom panel of Figure 8 ) amounts to a very significant ∼0.12 mag. Assuming a Gaussian distribution of magnitude errors, the faintest among our candidate AGB stars could actually be TRGB stars (or real AGB stars slightly fainter than the TRGB) which entered into the AGB selection region just because of the photometric scatter. However, we verified that this fraction is very small: just 55 AGB candidates (shown in Figure 8 ), or ∼6 % of the sample, have a F160W magnitude compatible with it being a TRGB star scattered upwards by one rms. These cases include AGB candidates in dense globular clusters such as AP 1969 (B225-G280; Figure 4 ). Overall, these estimates of the cluster densities and crowding errors are quite crude and subject to many uncertainties: few RGB stars per cluster, no real fitting of cluster density profile, no real estimate of the likelihood of RGB stars being scattered into the AGB region, etc. Moreover, it suffices to examine Figure 4 to conclude that proper estimates of the effect of crowding would require extensive—and time-consuming—experiments of artificial star tests using the actual cluster images, or fully synthetic cluster tests. For the moment, we limit ourselves to making our estimates available in the final catalog, so that potential users can appreciate, with the help of Figure 8 , the chances that AGB candidates in a particular cluster are significantly affected by crowding. 2.6. C/M Classification with WFC3/IR Medium-band Filters We augment our catalog of cluster AGB candidates with estimates of their spectral subtypes where available. First, we use results from Boyer et al. ( 2013 , 2019 ), who performed an HST survey using WFC3/IR medium-band filters in 21 regions across the PHAT footprint, to classify their bright red giants as either C- or M-type. Their method uses the sensitivity of the filters F127M, F139M, and F153M to the presence of either CN + C 2 or H 2 O absorption features in the NIR. We cross-match our sample of 937 candidate AGB stars in clusters with the complete Boyer et al. ( 2019 ) catalog. We find a total of 59 robust matches, for which the differences between the catalogs are smaller than 004 in position, and smaller than 0.017 mag in both F110W and F160W magnitudes. All of them are classified as M stars. The lack of C stars is this small sample is not surprising, given the remarkable scarcity of C stars in M31, especially in its innermost regions (Boyer et al. 2019 ). Indeed the average C/M ratio across the M31 disk amounts to just ∼0.02. 2.7.
Paper: 2018MNRAS.474.5437L (acknowledgments)
Acknowledgements CL was supported by the ILP LABEX (under reference ANR-10-LABX-63 and ANR-11-IDEX-0004-02) during the first part of this work and is now supported by a Beecroft Fellowship. JD and AS acknowledges funding support from Adrian Beecroft, the Oxford Martin School, and the Science and Technology Facilities Council (STFC). OI acknowledges the funding of the French Agence Nationale de la Recherche for the project ‘SAGACE’. This work relied on the HPC resources of CINES (Jade) under the allocation 2013047012 and c2014047012 made by GENCI and on the Horizon Cluster hosted by Institut d'Astrophysique de Paris. We warmly thank S. Rouberol for running the cluster on which the simulation was post-processed. This research is part of Spin(e) (ANR-13-BS05-0005, http://cosmicorigin.org ) and horizon-UK. This research is also partly supported by the Centre National d'Etudes Spatiales (CNES). This work is based on data products from observations made with ESO Telescopes at the La Silla Paranal Observatory under ESO programme ID 179.A-2005 and on data products produced by TERAPIX and the Cambridge Astronomy Survey Unit on behalf of the UltraVISTA consortium. We thank D. Munro for freely distributing his Yorick programming language and OpenGL interface (available at http://yorick.sourceforge.net/ ).
Paper: 2016ApJ...821...21C (fulltext)
; Kudoh Basu 2003 ; Wiebe Watson 2004 ; Falceta-Gonçalves et al. 2008 ; Hildebrand et al. 2009 ; Houde et al. 2009 ). Reduction of due to averaging effects is of great importance in our paper. Roughly speaking, two main averaging effects exist. First, averaging along the LOS can reduce . That is, if there is more than one independent turbulent eddy along the LOS, the measured value of will be reduced (see Myers Goodman 1991 ; Zweibel 1996 ; Houde et al. 2009 ). Second, averaging the polarization angle within the telescope beam can also reduce (see Heitsch et al. 2001 ; Wiebe Watson 2004 ; Falceta-Gonçalves et al. 2008 ; Houde et al. 2009 ). The averaging effects in general make the step from Equation ( 2 ) to Equation ( 3 ) inaccurate. In the presence of the averaging effects, the CF method tends to overestimate . In this paper, we focus on the averaging effect along the LOS and propose a simple method to compensate for the effect. Using three-dimensional direct magnetohydrodynamic (MHD) turbulence simulations, we test the proposed technique. In Section 2 , we describe our numerical method and explain in detail how averaging along the LOS makes the conventional CF method overestimate . At the end of Section 2 , we describe our new technique. In Section 3 , we present the results of our numerical simulations. In Section 4 , we discuss our findings and summarize. 2. NUMERICAL METHOD AND THEORETICAL CONSIDERATIONS 2.1. Numerical Code To obtain turbulence data, we solve the following compressible MHD equations in a periodic box of size using an essentially non-oscillatory scheme (see Cho Lazarian 2002 ): with ∇ · and an isothermal equation of state , where C s is the sound speed and ρ is density. Here, is the velocity, is the magnetic field, and is the driving force. We use 512 3 grid points. In our simulations, C s = 0.1, , and . In all simulations, the rms velocity v rms is between ∼0.7 and ∼0.8, and the sonic Mach number is . Since the Alfvén speed of the mean field () is 1, the Alfvén Mach number is , which means that the turbulence considered in this paper is sub-Alfvénic. 2.2. Forcing In this work, we drive turbulence in Fourier space and consider only solenoidal () forcing. We use ∼100 forcing components isotropically distributed in the range , where k is the wavenumber and k f varies from simulation to simulation (see Table 1 ). Therefore, the peak of energy injection occurs at . More detailed descriptions of forcing can be found in Yoo Cho ( 2014 ). Table 1 Simulations Notes. a Sonic Mach number. b Alfvén speed of mean magnetic field. c Driving wavenumber. d L los is the system size and L f is the driving scale. In our simulations, L los = 2 π and . Therefore, we have . 2.3. Theoretical Considerations: The Effects of Averaging along an LOS Our main concern is to investigate the effect of the driving scale on the CF method. For this purpose, we change the driving scale by changing the driving wavenumber k f . If the driving scale is L f ,
Paper: 2020A&A...635A.174V (fulltext)
region (black dashed line). For the viscous simulation on the other hand, we can see that the Poynting flux leaving the coronal shell region (solid black line) is substantially smaller than the Poynting flux entering the region (green dashed line). This difference is of the same order as the internal energy increase in the shell region (not shown here). The Poynting flux entering the shell region through the near TR-coronal boundary (green lines) differs in the ideal and viscous simulation due to the different mixture of propagating and standing waves in this region. For the ideal simulation, the standing wave component dominates, which reduces the amount of Poynting flux through the near boundary, as the standing regime creates a node in v z (see e.g. Prokopyszyn et al. 2019 ). Comparing with coronal energy requirements, the difference between the average Poynting flux entering and leaving through the TR-coronal boundaries of the shell region for the viscous simulation is of the order of 1.5−2 Jm −2 s −1 , several orders of magnitude too low, even for the Quiet Sun (∼3 × 10 2 Jm −2 s −1 , Withbroe Noyes 1977 ). Fig. 9. Time integrated Poynting flux , with S the boundaries of the left shell, for the viscous (solid lines) and ideal simulation (dashed lines). The green lines represent the lower TR-coronal boundary and the black lines the upper boundary. 4. Discussion and conclusions In this paper, we use 2.5D numerical simulations to investigate whether a pre-existing density profile is modified by the evaporative upflows following (viscous) heating by phase mixing of Alfvén waves. The Alfvén waves are driven through an additional forcing term in the momentum equation. The simulations account for the effects of gravitational stratification, thermal conduction, and optically thin radiation. A fully resolved atmospheric model (including the chromosphere as a mass reservoir) is achieved with relatively modest resolution by artificially broadening the TR using the approach proposed by Lionello et al. (2009) and Mikić et al. (2013) . To include a transverse density profile (perpendicular to the background magnetic field) in the initial setup, a heating function which varies in the cross-field direction is imposed. Throughout the simulations (running for a total time of more than 5000 s), a complex combination of Alfvén waves ( v z ) and longitudinal perturbations ( v y ) is present in the domain. The longitudinal perturbations are a combination of long-period oscillations resulting from the ongoing evolution due to the imposed cross-field heating profile, the ponderomotive effects associated with the driven Alfvén waves, and the evaporative upflows resulting from heating occurring in the coronal part of the shell regions of the loop due to phase mixing of the Alfvén waves. By comparing with the results from the ideal simulations, we are able to distinguish the evaporative upflows present in the viscous simulations from the other (ideal) longitudinal
Paper: 2016A&A...589A...9C (fulltext)
PDSs obtained during the 20042005 outburst decay ( Qu et al. 2005 ). They found a QPO at 0.22 Hz in addition to a narrow peak (i.e. pulsation from the pulsar) close to that frequency (0.23 Hz). The QPO was detected when the luminosity was at a mid-flux level of the outburst, or even lower. The rms strength of this QPO also amounts to a few per cent in that work. The 0.22 Hz QPO was observed around the spin frequency in both the INTEGRAL ( Mowlavi et al. 2006 ) and RXTE data during the 2005 outburst decline ( Qu et al. 2005 ; Reig et al. 2006 ). In addition, Takeshima et al. (1994) discovered another QPO with a centroid frequency of 0.05 Hz and a relative rms amplitude of 5% in the 2.337.2 keV data from the 1989 outburst of V 0332+53. We do not find any signal of the QPO at 0.05 Hz, although we find the presence of a broad component at around that frequency (which may eventually turn into peak noise, i.e. QPO). The QPO detected in V 0332+53 at 0.22 Hz constitutes the first detection of a QPO riding on the spin frequency of a NS and points to a strong coupling between the periodic (i.e. the pulses) and the red-noise (i.e. broad-band QPO signal) components ( Qu et al. 2005 ). The existence of a QPO signal with a similar frequency to the spin of the NS suggests that they are related in some way. One model to explain this coincidence was given by Lazzati Stella (1997) and Burderi et al. (1997) . They proposed that the QPO originates in the magnetosphere around the NS. Additionally, they argued that the accretion flow is inhomogeneous near the surface of the NS, and that random shots are characterized by an arbitrary degree of modulation. The combination of rotation and radiative transfer effects should therefore produce a periodic modulation of the shots similar to that of any continuum X-ray emission from the polar caps. They proposed that a coupling between the periodic and red-noise variability should be frequently present in X-ray pulsars. In our case and in contrast to previous observations during giant flares ( Reig et al. 2006 ; Reig 2008 ) the QPO signal appears during the lowest of the measured fluxes ( ). If the QPO has an origin on the magnetosphere of the NS (and is not an artefact from our observations), it should be more visible when the accretion column is more tenuous (or even absent). In summary, it looks as if we might be witnessing the innermost parts closest to the magnetosphere surroundings during our latest observations, i.e. Obs. 611. 5. Summary and conclusions In this paper we present a multiwavelength study of the Be/X-ray binary system V 0332+53 with the main goal of studying its transient behaviour during its intermediate luminosity X-ray outburst in 2008. After being yr in quiescence in X-rays, the system underwent a new active period. A new X-ray outburst was detected around 17 October 2008, roughly six days before the periastron passage of the NS. We note that the optical companion entered the brightening state nearly