adsabs/WIESP2022-NER huggingface dataset:
Paper: 2017ApJ...846...48D (fulltext)
K are highlighted. The inset summarizes the distribution of all shown planets over temperature. Overlaid are two lines of constant temperature. The 1000 K line corresponds to a change in aggregation observed in our work (see below). 1500 K is approximately the upper limit of the condensation region for typical silicates. The database is still one with low statistics, and detailed dependencies on temperature might be discussed. There is, e.g., a maximum at about 500 K, but the distribution might just as well be considered flat in a wide range. None of these statements are statistically significant at the moment. Nevertheless, there is a clear gap in the occurrence of planets between about 1000 K and 1500 K. For a planet forming at a larger distance from the star, the inward drift would not need to stop at the 1000 K line. Also, from a more general perspective of late planet formation phases there is no obvious reason why planets should not form, e.g., at 1100 K as likely as at 900 K. That leaves the question whether this temperature gap is inherited from the early phases of planet formation by aggregation at elevated temperatures. More often than not, planet formation and its limits are considered in light of radial distance. It might be argued that the lack of planets at higher temperatures might be due to a closer distance and a related lack of material for planet formation. Certainly, temperature and distance are related, but we note that they are not the same. Small changes in stellar mass already lead to large luminosity and significant temperature variations. However, the local material available will not significantly change. Also, it is currently often argued that inner planetary systems are formed from material drifting inward (Boley Ford 2013 ; Wurm et al. 2013 ; Hu et al. 2014 ). Inward drift is a classical mechanism for redistribution of matter (Weidenschilling 1977 ). As a special aspect, e.g., Ebel Alexander ( 2011 ) suggested that Mercury might have formed from dusty matter (chondritic porous interplanetary particles) drifting inward that were reprocessed in a region of higher temperatures with a composition that is then very different from solar composition. Considering such work, one motivation for our experiments is a radial inward drift of a mix of materials (see below). The mass of material available for planet formation does not necessarily depend only on the radial distance to the star. We also note that current planet temperatures and temperatures at the time of dust aggregation might not be the same. They might be comparable if the relevant aggregation occurred in a passive disk. They might have been higher in a still actively accreting disk. However, accretion of dust at these active times might have removed these aggregates again. In any case, we do not claim to have the only answer to planet formation in the inner disk at high temperature, but we think it is curious that there is a temperature limit in the exoplanet
Paper: 2016MNRAS.462..740B (fulltext)
to have natural transport from one planet to the next one in the sequence of bicircular problems, and a particle could drift away from one planet to reach a neighbourhood of the following planet. After that, for transport between neighbourhoods of the libration points of the same Sun–planet problem, it is enough to consider an RTBP (Sun+planet+infinitesimal particle), in which case the existence of heteroclinic connections between libration point orbits around L 1 and L 2 is well known. These connections would allow a particle to continue its journey towards the innermost Solar system. See Fig. 5 . Figure 5. Scheme of the chain of connections between two consecutive restricted bicircular Sun–Jupiter–planet (BR4BP ip ) problems, and two consecutive restricted Sun–planet ( P ip ) problems (RTBP ip ). For the computation of the connections, we proceed as follows. Consider two consecutive bicircular problems BR4BP ip and BR4BP ip + 1 corresponding to the planets ip and ip + 1. We are interested in transits from the outer to the inner Solar system. In all the bicircular problems, OL 1 is an inner orbit than the orbit of the planet with respect to the Sun, and OL 2 is an outer orbit (see Fig. 4 ). Then, the suitable connections are those involving the invariant manifolds of and . It is well known that the invariant manifolds associated with the equilibrium points L 1 and L 2 in the RTBP have two branches: one goes inwards, while the other one goes outwards, at least for times not too big. The invariant manifolds of the dynamical substitutes have the same behaviour. According to this, we denote by the branch of the invariant manifold that goes outwards, and by the branch of the invariant manifold that goes inwards. See Fig. 6 . Figure 6. Projection in configuration space (rotating coordinates) of branches of the invariant manifold (outer corona) and (inner corona). The dynamical substitute is the periodic orbit between the two branches. See the text for further details and also Fig. 8 . Therefore, in order to find connections, in the BR4BP ip , we compute the (linear approximation of the) parametrization of the stable manifold (see equation 7 ) and we follow the branch , and in the BR4BP ip + 1 , we compute the (linear approximation of the) parametrization of the unstable manifold (again see equation 7 ) and we follow the branch . To study if there exist connections, we also fix a section Σ R = {( x , y ); x 2 + y 2 = R 2 }, where R is an intermediate value between the radii of the orbits of the planets of the two bicircular problems, a ip and a ip + 1 , that is a ip R a ip + 1 (see Fig. 7 ). Figure 7. Projection in configuration space (rotating coordinates) of the dynamical substitutes , , one orbit on the invariant manifolds and , and the section Σ R for . The main objectives are: first, to determine if both manifolds reach the section Σ, and, secondly, to study which is the minimum distance between the sets and . That is, we want to see if both
Paper: 2017ApJ...849...98Z (fulltext)
small orbital eccentricities () of the terrestrial planets. Here, we propose that the observed wide gaps in some debris disks may have been cleared by the secular resonances of a known planetary candidate (in Fomalhaut) or yet to be detected (around Vega and other systems) isolated or multiple gas giants or super Neptunes, as the secular resonances sweep across this region during the depletion of the gas disk. Analogous to the asteroids, excited eccentricities of the residual planetesimals in the region between the inner disk and outer ring are mainly damped by the hydrodynamic drag of the diminishing disk gas. In addition to the extended gap, this scenario offers a potentially exciting prospect of finding signatures of ongoing planetary assemblage in debris disks around stars with an age comparable to the formation epoch of terrestrial planets in the solar system (Su et al. 2009 ). Since not a single planet with a mass in excess of (an observationally inferred upper limit) has been found around Vega and Fomalhaut, we explore here the possibility that the extensive gaps are produced by a sweeping secular resonance of a lone-planet with a modest mass and eccentricity in the debris disks around these and other similar systems. Although systems of multiple Neptune-mass planets may also lead to a similar effect, we opt the lone-planet scenario as an idealized simplest possibility. This paper is organized as follows. In Section 2 , we briefly introduce the lone-planet possibility in the Vega system, including detailed physical models to mimic the dust-free gap opening process by the lone-planet scenario in the Vega system, develop a rough method to constrain planet candidates based on observed features, and provide predictions for future planet detection. In Section 3 , we further extend our model to Vega-like system, Fomalhaut, inferring the properties of the prospective planet and compare these with a planet candidate, Fomalhaut b. A summary of our results and an extensive discussion on the other two-belt debris disks are presented in Section 4 . 2. A Lone-planet Scenario for the Vega System 2.1. Dynamical Causes for Wide Dust-free Gaps in Debris Disks Infrared observations indicate that Vega’s debris disk contains a large gap that spans from ∼15 au to ∼110 au (Su et al. 2013 ). The relatively large size of the gap in Vega’s debris disk can be used to constrain the orbital configuration of Vega’s (hypothetical) planetary system. In our solar system, the four gas giants cleared most of the debris between the outer asteroid belt ∼3 au and the inner Kuiper Belt ∼35 au. Similarly, in the young HR 8799 system, (at least) four giant planets have been detected (Marois et al. 2008 , 2010 ) and they are believed to be the chief culprits for the large separation between the warm (∼6–15 au) and cold (∼90–300 au) debris disks (Su et al. 2009 ). Therefore, a multiple-planet configuration is a preferred scenario in the Vega and Vega-like systems, which has a
Paper: 2018ApJS..239...18A (fulltext)
during DES observations (for details, see Burke et al. 2017 ). For example, the air mass adopted for the standard atmosphere, 1.2, is well matched to the median air mass of exposures entering the co-add, 1.22. The system response is defined in steps of 5 from 3800 to 11000 . Out-of-band light leakage has been directly measured with DECal to be ≲10 −3 relative to the in-band response, and vendor measurements of witness samples suggest that the out-of-band leakage is typically at the 10 −5 to 10 −4 level. While detailed characterization of the out-of-band response is ongoing, the throughput of the DES DR1 standard bandpasses is defined as zero for out-of-band wavelengths (caveats are mentioned in Section 4.6 ). 5.5. Software All software used in the DESDM pipelines described in Section 3 and in Morganson et al. ( 2018 ) can be accessed from the release page itself 91 91 https://des.ncsa.illinois.edu/releases/dr1/dr1-docs/processing or from the DES Github Organization. 92 92 https://github.com/DarkEnergySurvey Together with access to the software used to generate these products and the configuration described in Morganson et al. ( 2018 ), we provide the main ingredients to reprocess the data in a manner similar to that done by DES. 6. Data Access Access to the DES DR1 data is provided through a collaborative partnership between NCSA, 93 93 http://www.ncsa.illinois.edu/ LIneA, 94 94 http://www.linea.gov.br/ and NOAO. 95 95 https://www.noao.edu/ From these institutions, a rich and complementary set of tools and interfaces were developed to access and interact with DES data in different ways that cover a broad set of use cases that enable scientific discovery. In this section we describe the main tools provided to access DES DR1. 96 96 https://des.ncsa.illinois.edu/releases/dr1/dr1-access 6.1. NCSA DES access NCSA provides the primary set of basic web applications to access DES DR1 data. Developed at NCSA, DESaccess ( https://des.ncsa.illinois.edu/easyweb ) provides the user with an interface to submit asynchronized jobs to perform SQL queries against an Oracle DB that contains the DES DR1 catalogs and to generate cutouts from a given list of positions from the co-add images. It also contains information regarding the catalogs, an interactive footprint finder to locate positions and tiles, and means to access images and catalogs in a tile-based format. The main services can be summarized as follows: 1. SQL web client: We provide an SQL web client that allows the user to submit asynchronous query jobs against the Oracle 12 database that contains the DES DR1 tables. The submitted jobs enter into a queue, and results can be retrieved at later times in either csv , FITS (Wells et al. 1981 ), or HDF5 (The HDF Group 1997 ) file format, supporting compression in some of the cases. Big jobs are divided into chunks of 1.5 GB to facilitate file transfer. This interface also provides means to check syntax and to evaluate synchronous jobs for a quick
Paper: 2021ApJ...911...61B (fulltext)
by larger quiet regions in which B varied relatively slowly, and contained small amplitude waves and turbulence. Thus, the basic structure of the VLISM observed between 2012/DOY 238 and 2020/DOY 182 consists of four jump–ramp structures associated with the four maxima in B between 2012 and 2020, separated by four quiet intervals. The timescales of the jumps that moved past V1 were ∼5.4 days and 3.3 days for the shocks sh1 and sh2 observed on ∼2012.9746 (DOY 335) and 2014.6 (DOY 236), respectively. The corresponding timescales for the jumps associated with pressure fronts were ∼35 days and ∼8 days for the pressure fronts pf1 and pf2 observed on ∼2016/346 and ∼2020/147, respectively. The corresponding timescale for the ramp associated with sh1 was ∼162 days, during which time B decreased exponentially from ∼2012.9746/335 to ∼2013.35759/131. The timescale of the ramp associated with sh2 was 276 days, when B decreased linearly (with notable fluctuations related to fine-scale features) from 2014.68212 to 2015.37669. In the ramp following the pressure front pf1, B declined nearly monotonically until 2018.41885, during an interval of 402 days. We do not yet have V1 observations during the ramp associated with pf2. We have shown that jump–ramp features of B in the VLISM and the larger quiet intervals between them are the fundamental features of the VLISM. These features were actually specially configured waves that were generated by the interaction of heliosheath with the heliopause and then propagated through the VLISM, presumably interacting with one another to produce the results observed by Voyager 1. Smaller scale waves and turbulence exist within these features. Mathematical techniques such as Hilbert transforms and wavelet transforms (Zhao et al. 2020 ) and Fourier methods (Fraternale et al. 2019 ) are very valuable tools for studying the VLISM over a wide range of scales. However, it is also essential to understand the basic structure and dynamics of the large-scale flows on scales of hours to a decade presented in this paper. This can best be done by improving the MHD models and the corresponding input observations. A number of models have been discussed in the literature, including those by Fermo et al. ( 2015 ), Zirnstein et al. ( 2018 ), and Kim et al. ( 2017 ). However, none of these models reproduce the kind of pattern of shocks, pressure pulses, and quiet regions shown in Figure 2 . Appendix New Method for Deriving Plasma Densities The Voyager PWS instrument (Scarf Gurnett 1977 ) includes a wideband waveform receiver that captures digitized waveforms in the range of 50 Hz to above 10 kHz. Waveform captures include 1600 4 bit samples obtained at a rate of 28,800 samples per second. On board, 800 of these waveform captures are obtained at a rate of one every 60 ms covering 48 seconds (one frame). Due to decreasing telemetry performance at the large distances of Voyager, beginning in 1992 November, only one of every five waveform
Paper: 2019MNRAS.484.5845L (fulltext)
are computed by integrating the hydrostatic equations with the equation of state given by where p is the pressure, ρ is the mass density, G is the gravitational constant, p b is the pressure at the base of the radiative envelope, and and a is the isothermal sound velocity, and R J is the radius of Jupiter. In the convective core where p ≫ p b , we obtain corresponding to a polytrope of the index n = 1, and hence for Γ 1 ≡ (∂ln p /∂ln ρ) S ≈ 2 the square of the Brunt–Väisälä frequency where On the other hand, in the radiative envelope where p ≪ p b , we have and and c denotes the adiabatic sound speed. Since p /ρ ≈ a 2 ∝ T for an ideal gas, the temperature T is constant for a constant a , suggesting an isothermal atmosphere. In this paper, following Auclair-Desrotour Leconte ( 2018 ), for the Jovian planet model of mass M = 0.7 M J with M J being the Jovian mass, we use p b = 100 bar = 10 8 dyn cm −2 and set the outer boundary R e at p = 0.01dyn cm −2 and define the planet’s radius R = R e /1.01 = 9.31 × 10 9 cm. We also assume Γ 1 = 2 in the interior. Fig. 1 is the propagation diagram where N 2 and with l = 2 are plotted as a function of r / R where L l is the Lamb frequency. Figure 1. Squared characteristic frequencies N 2 (solid lines) and (dashed lines) plotted as a function of r / R (left-hand panel) and of log p (right-hand panel), where N and L l are, respectively, the Brunt–Väisälä frequency and Lamb frequency and δ A = 10 −4 is assumed for N in the convective core. Note that the squared frequencies are normalized by GM / R 3 where M and R are the mass and radius of the planet, G is the gravitational constant and l = 2 is used for L l . The dotted line plots log 10 | N 2 | in the convective core where N 2 0. 2.2 Perturbed basic equations The perturbed basic equations for a uniformly rotating planet may be given by where ρ, p, s, T , and c p are, respectively, the mass density, pressure, specific entropy, temperature, and specific heat at constant pressure, is the energy flux vector, is the displacement vector, α T = −(∂ln ρ/∂ln T ) P , and (′) and δ indicate, respectively, the Eulerian and Lagrangian perturbation of quantities. Note that we have applied the Cowling approximation neglecting the Euler perturbation of the gravitational potential due to self-gravity, and that we have ignored the effects of rotational deformation. The vector denotes the angular velocity vector of rotation and is assumed to be constant for uniformly rotating planets. Here, we have assumed that the perturbations are proportional to the factor e iω t + i m φ , where ω = σ + m Ω with is the oscillation frequency in the co-rotating frame of the planet and σ is the frequency observed in the inertial frame, and m denotes the azimuthal wavenumber around the axis of rotation. Since evolutionary models of irradiated Jovian planets computed with appropriate equations of state and opacities are not available to us, we employ an approximate treatment to estimate the
Paper: 2017ApJ...843L..13Z (fulltext)
be directly detected. In this case, T may be adopted as a small value . Our observations require that the maximum γ -ray luminosity at the relevant age T is smaller than the upper limits, i.e.,where , and is the eight-year limit presented in the last row in Table 1 . Given a set of magnetar parameters P 0 and B p , the spin-down luminosity depends on the age of the magnetar, T . Since the earliest burst, FRB 121102, was detected in late 2012, the magnetar needs to be at least four years old until now. We consider two possibilities. First, if the pulsar was born between 2009 January 1 and 2012 November 2, one should use the data during that period of time to constrain magnetar parameters assuming different transparent times by setting in Equation ( 1 ). The transparent time depends on the amount of ejecta and opacity, which may range from minutes (for NS–NS mergers) to ∼ years (for supernovae). In our calculation, we use 10 2 s, 10 4 s, 0.1 years, 0.5 years, and 1 year, respectively. Since the upper limits are slightly different in different time bins, we use the average upper limits to derive the constraints as presented in Figure 2 . Figure 2. Constraints of the pulsar parameter by the LAT upper limit. Solid lines: cases of dotted lines: cases of dashed lines: cases of . The horizontal line denotes the break-up limit P 0 = 0.6 ms. The parameters are constrained to the regions that are above the curves/lines. The second possibility is that the magnetar was born before 2009 January 1. For these cases, the time T should be adopted as the age of the magnetar when the first observation started (in 2009 January 1). When the birth time is much earlier than 2009 January 1, one can use the tighter upper limit over a much longer time (e.g., four-year or eight-year) limit to perform the constraint. For this case, we consider T = 10 years as an example. The results for , 10 4 s, 0.1 years, 0.5 years, 1.0 year, and 10 years are plotted in the plane ( and assumed) as shown in Figure 2 . For each T , we consider , , and . The region above each curve is the parameter space allowed for the given T and η value. Since the minimum spin period of a millisecond pulsar would be the break-up limit ms (e.g., Vink Kuiper 2006 ), our results suggest that there is essentially no constraint from the upper limits if the pulsar age in the case of . This means that the current limit cannot constrain the magnetar parameter if it was born before Fermi observations started. On the other hand, for and if the magnetar age is 4–8 years and the transparent time shorter than 0.1 years, the upper limits in the first four years could place some constraints on B p of the pulsar. In particular, if the transparent time is short enough, the non-detections exclude high magnetic field strengths. As a result, if FRB 121102 progenitor was indeed born after 2009 January 1 and has a magnetar origin, it must have been born in a dirty environment with a delayed transparent time (say years) in
Paper: 2016MNRAS.459.1533V (fulltext)
magnetic fields across the solar surface (Mackay Yeates 2012 ). Although synoptic maps of the stellar vector field exist for a couple of decades and have been derived for hundreds of systems (e.g. Donati Collier Cameron 1997 ; Donati et al. 1999 , 2006b ; Donati Landstreet 2009 ), there are significantly fewer synoptic maps of the vector field of the Sun. Earlier synoptic maps of the vector solar field were created on the basis of the rotational modulation of LOS magnetograms (Pevtsov Latushko 2000 ; Wang Zhang 2010 , also called ‘pseudo-vector reconstruction method’). With this method, these authors assumed that the field on large spatial scales does not change over several days. Because of this assumption, these early synoptic maps reconstructed only large-scale magnetic fields of the Sun. More recently, Gosain et al. ( 2013 ) used daily full-disc vector magnetograms from the SOLIS/Vector Spectromagnetograph (VSM) spectrograph to reconstruct the vector solar field of a series of 23 Carrington Rotations. This is a significant advance compared to the large majority of synoptic charts of the solar magnetic field, which are restricted to the LOS component of the solar magnetic field, used to derive B r . An important conclusion reached by Gosain et al. ( 2013 ) is that the solar magnetic field is not normal (i.e. purely radial) to the surface at the photospheric level. This is somewhat disturbing, as the assumption that the solar magnetic fields are purely radial are used in several studies. With the new existing capabilities in producing vector synoptic fields, we can now start to investigate the behaviour of the meridional B θ and azimuthal B φ components of the solar magnetic field and investigate how the three magnetic field components in the Sun compare to the ones derived in stellar observations (e.g. Fig. 1 ). However, direct comparison between stellar and solar magnetic field topologies are hampered by their dramatic differences in resolutions. To provide a direct comparison between solar and stellar magnetic field synoptic maps, it is first required to ‘filter out’ the small-scale field (i.e. the finely resolved magnetic features) of the solar observations, since this component is currently not accessible to ZDI studies. After this, we are left with only the solar large-scale field, whose spatial scale is comparable to those achieved by stellar observations. One way to filter out the small-scale structure of the solar field is to express the solar magnetic field as spherical harmonics of different l degrees, up to a maximum degree l max . The smallest l degrees represent the largest-scale components, e.g. l = 1 for the dipole, l = 2 for the quadrupole, l = 3 for the octupole and so on. By retaining only the components of low l degrees in the spherical harmonics decomposition, we are then able to calculate the large-scale component of the magnetic map. This approach has been extensively applied to the LOS and radial components of the