adsabs/WIESP2022-NER huggingface dataset:
Paper: 2017A&A...603A..96R (fulltext)
indeed become an important ionization source in T Tauri disks. However, our model should be seen as a first attempt at a comprehensive modelling of SP ionization in protoplanetary disks. In the following we discuss further important aspects like variability, non-stellar origin of SPs, the importance of magnetic fields and future prospects for SP modelling in protoplanetary disks. 4.3.1. Flares and variability In our models we assumed continuous (i.e. time-averaged) particle fluxes and X-ray luminosities. Although this is a reasonable assumption (see Sect. 2.2 ) it is likely that the disk is also hit by singular powerful X-ray and/or SP flares. From the X-ray COUP survey of the ONC cloud, Feigelson et al. (2002) and Wolk et al. (2005) derived a median X-ray flare luminosity of , comparable to the X-ray luminosity in our high X-ray models, and peak flare luminosities up to . The duration of such flares can last from hours up to three days with a typical frequency of roughly one powerful flare per week. Ilgner Nelson (2006) argue that for such flare properties the disk ion chemistry responds to time-averaged ionization rates. However, in case the duration between two flares is longer than the recombination time scale in the disk, such singular and strong flares would produce asymmetric features in molecular ion emission (i.e. a singular flare only affects a certain fraction of the disk). The spatially resolved SMA observations of TW Hya show indeed such an asymmetric structure ( Cleeves et al. 2015 ). However, as discussed by Cleeves et al. (2015) these features could also have a different origin like spiral arms or an hidden planet locally heating the disk. If such asymmetric structures are caused by stellar flares they would provide complementary constraints on the X-ray/SP activity of the star. Multi-epoch data of spatially resolved molecular ion emission is required to prove the flare scenario (i.e. the features should disappear quickly). The current (sub)mm interferometers like ALMA, NOEMA and SMA provide the required spatial resolution and might even allow for monitoring of disks in molecular ion lines on a daily/weekly basis in the future. Modelling of such observations does not necessarily require 3D chemical models. Assuming that the disk physical structure is not affected by flares, radial cuts through the disk can be modelled with 2D (time-dependent) thermo-chemical models as presented here. 4.3.2. Magnetic fields In our model we neglect the impact of magnetic fields on the SP transport. Stellar and disk magnetic fields are in particular relevant for the question if particles actually hit the disk (see Feigelson et al. 2002 , for a discussion). Magnetic fields can either drag the particles away from the disk (e.g. like in a wind) but could also funnel the particles and concentrate the ionizing flux in particular regions of the disk. In the first scenario SPs can still have an impact on the upper layers of the disk but certainly not on
Paper: 2018MNRAS.476.4162O (fulltext)
have a constant absolute magnitude of −7 mag. The investigation on decline times for all novae in the catalogue indicates that the novae with D- and J- type light-curve shapes may not have linear decline. If so, only considering the decline times ( t 2 or t 3 ) may not characterize the absolute magnitude at outburst maximum and the decline well. All of these arise uncertainties when using of MMRD relation if the type of nova is not clarified. Furthermore, the luminosities of novae may differ depending on their stellar populations or their evolutionary states, i.e. on the stellar age and metallicity. For the scatters found in other distance indicators, e.g. Cepheids, it has been suggested that a different metallicity of the host galaxy may significantly affect the zero-point of the relation (Freedman et al. 2001 ). In this case, there is a possibility on any environmental dependence on MMRD relations as well. Since our and the other studies in literature depend on disc novae in the Galaxy, the contribution of other populations (bulge, halo, even thick disc) is not well known. Our study does not dispel the concerns regarding the validity the MMRD relation, but has raised open questions. To understand the MMRD relation better, the novae in galaxies, which have reliable distances (and extinctions), should be studied with high quality photometric and spectroscopic data. A larger sample of them would also be very welcome as well as corresponding parameters and well-determined light curves that allow us to classify their light-curve shapes. Our catalogue may also help to investigate the nature of MMRD diagrams in the long-term. Studying the MMRD relation is important to understand the origin of the nova eruption itself, besides using it as a tool to estimate absolute magnitudes. In this study, we did not find any signs that the spectral classification represents the stellar population of the Galactic novae, as suggested by Della Valle Livio ( 1998 ). However, our decline time histogram implies that the He/N+Hybrid novae are fast systems with a peak time at t 3 = 13 d, while the slow Fe ii novae have a peak decline time at t 3 = 42 d. In contrast to the suggestion of Della Valle Livio ( 1998 ), we found that both Fe ii and He/N novae or fast and slow novae are located in the whole range of vertical distances, but mainly concentrated throughout the Galactic disc. The differences between the results in these two studies may arise from the samples used in the analyses. The sample in Della Valle Livio ( 1998 ) contains a smaller number of novae with spectral classification, which reach only to a vertical distance of ∼2 kpc, whereas our sample have much more systems reaching up to two times farther in vertical direction from the Galactic plane. A relatively large number of slow novae located close to the Galactic plane and fast novae farther from Galactic plane must have been undetected in their analyses. The approach proposed by Della Valle Livio ( 1998
Paper: 2020MNRAS.491.5925M (fulltext)
INT-G (intermediate gravity) and FLD-G (field gravity) classes defined by Allers Liu ( 2013 ). In Section 2 , we describe the selected sample of objects presented in this paper. In Section 3 , we provide details of the X-shooter observations and data reduction. In Section 4 , we explain the analysis carried out and our results. In Section 5 , we derive the physical parameters for our sample using evolutionary models for substellar objects. Finally, in Section 6 we summarize our results. 2 SAMPLE SELECTION We selected a sample of 20 benchmark objects with confirmed membership of open clusters or associations, with known ages, metallicities and distances. We chose a sample with photometry and optical spectral types confirming membership. In the case of the Pleiades and Praesepe, members are confirmed not only by photometry and optical spectroscopy, but also by proper motion (Lodieu, Deacon Hambly 2012b ; Boudreault et al. 2012 ; Bouy et al. 2015 ). We selected objects in each of the associations or open clusters with magnitude ranging from J = 14.9 to J = 18.1 mag, which allowed us to obtain spectroscopy with a signal-to-noise ratio of between 4 and 53 in the near-infrared, using exposure times of between 40 min and 1.5 h. Our sample has spectral types between M5.5 and L3.5, with at least one target per open cluster or association of each spectral type between M5.5 and L0. To complete our sample in age range, we added three members from the UppSco association, and four from the Cha I region from proposals: 093.C-0109(A) (P.I. Van der Plas), 093.C-0769(A) (P.I. Bonnefoy) and 095.C-0378(A) (P.I. Testi). We list the objects in our sample with their main characteristics in Table 1 . Table 1. List of observed targets with their full name, coordinates, magnitude, spectral types, cluster or association membership, and flag for disc-bearing sources and binary candidates. Notes . References: [1] Boudreault et al. ( 2010 ), [2] Hodgkin et al. ( 1999 ), [3] Boudreault et al. ( 2012 ), [4] Bouvier et al. ( 1998 ), [5] Stauffer et al. ( 1998 ), [6] Bouvier et al. ( 1998 ), [7] Bouvier et al. ( 1998 ), [8] Nagashima et al. ( 2003 ), [9] Lodieu, Deacon Hambly ( 2012a ), [10] Bihain et al. ( 2006 ), [11] Ardila, Martín Basri ( 2000 ), [12] Martín, Delfosse Guieu ( 2004 ), [13] Lodieu et al. ( 2007 ), [14] Persi et al. ( 2000 ), [15] López Martí et al. ( 2004 ). References of disks detections: [16] Luhman Muench ( 2008 ), [17] van der Plas et al. ( 2016 ), [18] Herczeg et al. ( 2009 ), [19] Long et al. ( 2017 ). a NR: the existence of a disc has not been reported for this object. References for spectral type: [20] Alves de Oliveira et al. ( 2012 ), [21] Luhman ( 2007 ), [22] Luhman ( 2004 ), [23] Lodieu et al. ( 2018 ), [24] Slesnick et al. ( 2008 ), [25] Bihain et al. ( 2010 ), [26] Martín et al. ( 1996 ), [27] Stauffer et al. ( 1998 ), [28] Boudreault Lodieu ( 2013 ), [29] West et al. ( 2011 ), [30] Manara et al. ( 2017 ). References for binary candidates:
Paper: 2018ApJ...853...22X (fulltext)
Σ after core collapse stops. The disk lifetime is close to that shown in Figure 6 . See the text for details. In Figure 8 , we show M disk and M * as functions of t for the disk evolutions ( ω = 0.3 × 10 −14 s −1 , T cd = 10 K, and M cd = 1.0 M ) with and without external FUV fields. The evolutions of M disk and M * are similar to those shown in Figure 3 . Since ω is smaller, more material of the core falls onto the star, which causes the star to grow faster and the disk to grow relatively slower. When the collapse ends, M * is about 0.86 M and M disk is about 0.14 M . Therefore, the gravitational potential well is deeper, which results in the weaker external FUV photoevaporation. Finally, the difference between the final star masses shown in Figure 8 is smaller (0.002 M ). Figure 8. Evolution of M disk (top panel) and M * (bottom panel) as functions of t for the same core properties as those shown in Figures 6 and 7 . The solid lines represent the case for including both internal X-ray and external FUV photoevaporations. The dashed lines represent the case for including internal X-ray photoevaporation only. Figure 9 shows R disk as a function of t for the disk evolutions ( ω = 0.3 × 10 −14 s −1 , T cd = 10 K, and M cd = 1.0 M ). The evolutions of R disk are similar to those shown in Figure 4 . However, the lower ω means that the location of core mass falling onto the disk is closer to the star. So, R disk for the disk without external FUV fields shown in Figure 9 is smaller than that shown in Figure 4 . As discussed in previous paragraphs, the star gets more mass during the core collapse for the case of ω = 0.3 × 10 −14 s −1 . Therefore, R g of the disk with external FUV photoevaporation grows quickly and it can finally reach ∼372 au. Although external FUV photoevaporation is very weak in this case, the gravitational potential well is shallow at the large radius of the disk. The external photoevaporation can still cut off the disk and decrease R disk effectively, as shown in Figure 9 . Figure 9. Evolution of R disk and R g as a function of t for the same core properties as those shown in Figures 6 and 7 . The solid line represents the evolution of R disk in the disk, including both internal X-ray and external FUV photoevaporations. The dashed line represents the evolution of R disk in the disk, including internal X-ray photoevaporation only. The dotted line represents the evolution of R g in the disk, including both internal X-ray and external FUV photoevaporations. Next, we discuss the evolution of external FUV photoevaporating disks with high ω . We first present the results of the disk evolution, including internal X-ray photoevaporation only. Figure 10 shows the evolution of Σ of the internal X-ray photoevaporating disk for the high low ω case that ω = 5.0 × 10 −14 s −1 , T cd = 10 K, and M cd = 1.0 M . In this case, the angular momentum of the pre-stellar core is very high. Most material of the core collapses onto the disk and is far
Paper: 2018MNRAS.473.1332G (fulltext)
in other directions form gas circulations in the meridional planes, and near the bottom of the AGN bubble, meridional vortexes also appear. At t 700Myr, one can clearly see that the trailing flow is falling back to the cluster centre (see the overplotted velocity vectors), suggesting that the outflow does not leave the cluster eventually, but instead it is confined to the cluster. Figure 8. Temporal evolution of the concentration distribution ( Z ) of the tracer fluid in a simulation where the initial value of Z at t 0 is 1 if r 100kpc and 0 if r 100kpc. Arrows over-plotted in each panel indicate directions of gas velocities. From top to bottom, the time of each panel is t 300, 400, 500, 700Myr. The top colour bar represents Z in linear scale. It is clear that trailing flows behind the AGN bubble form real outflows in the ICM, which fall back to the cluster centre at later times, e.g. at t 700Myr. Trailing outflows play an important role in suppressing or delaying the central cooling catastrophe. As the rarefaction wave propagates away to large radii, gas inflows resume in the inner core region, as can be clearly seen in Figs 3 and 8 and discussed in Section 3.2 . If there are no trailing outflows, the gas density in the cluster centre would rise rapidly and a central cooling catastrophe would then develop. Trailing outflows uplifted by AGN bubbles transport gas from the inner cluster core to outer regions, suppressing gas accumulation at the cluster centre resulting from inflows along other directions and thus averting the central cooling catastrophe, which is delayed until trailing outflows fall back to the cluster centre at later times (see Section 3.6 ). Observational evidence for trailing outflows has been found in some cool-core clusters where the spatial distributions of heavy elements are anisotropic and preferentially aligned with the large-scale radio and cavity axes (Kirkpatrick, McNamara Cavagnolo 2011 ; Kirkpatrick McNamara 2015 ), suggesting that hot metal-rich outflows are indeed lifted up by AGN jets. Hydrodynamic simulations following both the jet and metallicity evolution in galaxy clusters would be helpful to further explore this scenario (e.g. Barai etal. 2016 , which however does not investigate the angular distribution of metals). Along the axes of trailing outflows, gas density is relatively high, as can be seen in the bottom two rows of Fig. 3 , and cold gas may drop out due to local thermal instability, forming cold filaments found in some galaxy clusters (e.g. Mittal etal. 2012 ; Tremblay etal. 2015 ). 3.6 Emerging picture of cyclic cool cores What is the long-term evolution of cool cluster cores after powerful AGN feedback events? After the onset of the AGN jet event, the cool core in our simulation is quickly overheated by shock dissipation, and rarefied by the rarefaction wave which transports energy to large radii. As the bow shock weakens and propagates further away, the inner core region starts to contract again
Paper: 2019A&A...630A..43M (fulltext)
flow pattern is consistent across the diamagnetic cavity boundary in the CSE frame, we suggest that the flow pattern inside the cavity is determined by the flow outside, depending on the local plasma and magnetic field conditions. Fig. 7 Summary of flow pattern of the accelerated ions inside and outside the diamagnetic cavity in the YZ plane of ( a ) the CSEQ frame and ( b ) the CSE frame. The dark gray object at the center of each figure indicates the comet nucleus, and the surrounding light gray shaded area is the diamagnetic cavity. The view of these figures is from the Sun, and the four regions depending on the clock angle of the Rosetta position are also shown. In each region, a main tailward flow and its aberration component are shown by a cross and an arrow, respectively. Figure 8 summarizes the ion flow pattern inside and outside the diamagnetic cavity in the CSE frame. The aberration pattern of the accelerated ions in the - Z CSE direction near the XZ plane seen in Fig. 8 a suggests that they are mass-loaded by local cold cometary ions, which is similar to the mass-loading of the solar wind ( Behar et al. 2016a , b , 2017 ). The mass-loading of the accelerated cometary ions would result in an acceleration of the local cold cometary ions in the + Z CSE direction due to the conservation of momentum in the system, but we did not detect such a flow pattern from ICA data. This suggests that their energy is very low, which would be out of the energy range in this study (40 eV), or more likely, the density of the accelerated ions is much lower than the local cometary plasma, meaning that the earning momentum for the cold cometary ions is negligible. Fig. 8 Summary of the flow pattern of the accelerated cometary ions inside and outside the diamagnetic cavity in ( a ) the XZ plane and ( b ) the XY plane of the CSE frame. Near the XY plane of the CSE frame ( Fig. 8 b), the accelerated ions flow tailward with an anti-cometward aberration. This indicates that the accelerated ions are deflected around the diamagnetic cavity. The mechanism that produces this flow pattern is still unknown, but the accelerated ions may have been deflected when they flew across the upstream boundaries such as the ion-neutral collisionopause or the SWIC boundary. Another possibility is the aforementioned ambipolar electric field that points outward from the nucleus ( Odelstad et al. 2018 ; Bercic et al. 2018 ). This electric field may cause the tailward-flowing ions to be deflected around the diamagnetic cavity, which results in the anti-cometward aberration of the flow in the XY plane in the CSE frame. As seen in observations, the flux of these ions inside the cavity is nearly one order of magnitude lower than that outside the cavity. This suggests that the accelerated ions are only partly transferred from the outside into the inside of the diamagnetic cavity. There may be an energy-filtering effect in the cavity boundary. If the ion transmission into the cavity is
Paper: 2019A&A...629A.129M (fulltext)
(− z rec for the counter-jet) for which a recollimation shock is formed due to the jet and wind momentum fluxes balancing each other in the y -axis direction (perpendicular to the jet propagation). This condition can be approximately expressed as follows ( Bosch-Ramon Barkov 2016 ): where Ṗ w = Ṁ w v w is the wind thrust. For γ j = 1.2 we obtain z rec = 1.4 × 10 12 cm ≈ 0.47 a , whereas for γ j = 3, z rec = 3.1 × 10 11 cm ≈ 0.10 a . Some combinations of Ṗ w , Ṗ j and γ j yield non-physical (complex) values of z rec in Eq. ( 4 ), which means that for those sets of parameters no recollimation shock is formed, as the momentum flux balance is never reached (see Yoon et al. 2016 , for a non-relativistic simulation where this effect is studied). Electrons are injected at , since the jet displacement in the x and y coordinates is small at z = ± z rec (≲0.005 a in the worst case). The particles are assumed to follow a power-law energy distribution with spectral index −2, typical for acceleration in strong shocks via the Fermi I mechanism (e.g Drury 1983 ), with a minimum energy of MeV. An exponential cutoff is also assumed: where E ′ is the electron energy in the FF. The normalization of Q ′( E ′) is obtained from the condition that the total power available for non-thermal electrons in the FF is . Electrons are convected from along the jet path while they cool down via adiabatic, synchrotron and IC processes. For simplicity, we neglect particle acceleration beyond , even though weak shocks and turbulence may (re)accelerate particles further downstream in the jets. Analytical expressions are used for the energy losses in the FF (see Longair 1981 ; Khangulyan et al. 2014 for synchrotron and IC losses, respectively). The characteristic timescale for adiabatic cooling is , where R j ( l )∝ l is the jet radius. The magnetic pressure at the accelerator (located at ) is characterized in the FF as a fraction η B of the total jet energy density: where and r j, 0 = R j ( z rec ) are the magnetic field and the jet radius at , respectively. For the same , this fraction relates to the ratio of magnetic pressure to stellar photon energy density in the LF as , which is useful to compare the expected radiation outputs of synchrotron and IC processes. This yields for γ j = 1.2, and for γ j = 3. The jet magnetic field is assumed perpendicular to the flow in the FF, since at the scales considered in this work, far from the jet launching point, a dominant toroidal component for the magnetic field is expected (e.g., Pudritz et al. 2012 ). Given the adopted assumptions plus frozen in conditions, the magnetic field at each jet point can be computed as: The cutoff energy is the maximum energy that the electrons can achieve in the accelerator region in the FF, and it is obtained as follows: We compare the acceleration timescale with the cooling timescale , where Ė ′ is the cooling rate accounting for adiabatic, synchrotron and IC losses. In this work we set η acc = 0.1 as a
Paper: 2017MNRAS.471.4776G (fulltext)
activities during the period of 2012 August 7–September 4. SDO/AIA 211 and 193 Å observations show a series of coronal loops that formed at the heated region between the filament channel and the CH boundary before and after the partial eruptions (Fig. 4 ). Some loops are observed almost simultaneously with the filament plasma flow from its southern to the extreme northern footpoint on August 7 and some loops just before the onset of eruptions. Before and after the formation of these loops, several EUV bright points associated with the small magnetic dipoles appeared and disappeared quickly between the DR and the CH and very close to the filament channel and its southern footpoint. As observed in SDO/AIA 304 Å images (black arrows, Fig. 4 ), the loop formation was accompanied by the plasma motion from the filament channel to the western boundary of CH. Initially, just a few minutes before the eruptions, the filament plasma, related to a part of the central section not involved in filament eruption, moved from the central part of the filament channel, reached the western boundary of CH and induced the onset of CH fading (August 8, 07:15 ut; Fig. 4 ). Later, after the filament partial eruptions and CMEs, several additional loops formed and the evacuated filament plasma moved from the southern part of the filament channel, reached the southern side of CH and led to the disappearance of the dimming region and almost the whole CH (at 20:00 ut ). Furthermore, the analysis of SDO/HMI magnetograms after the partial filament eruptions and the disappearance of DR and CH show the magnetic diffusion at the north side of the erupted northern section of the filament. Figure 4. SDO/AIA 304 Å ( T = 5 × 10 4 K) and SDO/AIA 211 Å ( T = 2 × 10 6 K) images. CH boundaries are enhanced with white colour. Some loops between the filament (blue colour) and CH surroundings are enhanced with green colour. The black rectangle enclosed the CS and NS, as well as the surrounding region of the filament, related with the possible magnetic reconnection between the filament footpoints and the CH boundary. With the black arrows we indicate the plasma that is moving towards the CH. Additionally, in order to study the perturbed region between the filament channel and the CH boundary, we selected a sector (15° × 12°) from this region based on the SDO/AIA 193 Å images and divided it into 11 segments or cuts, parallel to the solar equator (Fig. 5 a). We have used the method, similar to that reported by McIntosh et al. ( 2011 ), Shen, Liu Su ( 2012 ), Verwichte et al. ( 2013 ) and Ma et al. ( 2014 ). We elaborated the time evolution map of each cut during 10 h, starting at 00:00 ut on August 8, which is the preceding period to the CH disappearance. Fig. 5 (b) shows one part of this map that reveals the CH evolution and its disappearance related to the arrived plasma. Generally, all segments of this map show the filament plasma disturbances and wavy behaviour of the longitudinally