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1.
More than 80 giant planets are known by mass and radius. Their interior structure in terms of core mass, number of layers, and composition however is still poorly known. An overview is presented about the core mass M core and envelope mass of metals M Z in Jupiter as predicted by various equations of state. It is argued that the uncertainty about the true H/He EOS in a pressure regime where the gravitational moments J 2 and J 4 are most sensitive, i.e. between 0.5 and 4 Mbar, is in part responsible for the broad range \(M_{\mathit{core}}=0{-}18\:M_{\oplus }\), \(M_{Z}=0{-}38\:M_{\oplus }\), and \(M_{\mathit{core}}+M_{Z}=14{-}38\:M_{\oplus }\) currently offered for Jupiter. We then compare the Jupiter models obtained when we only match J 2 with the range of solutions for the exoplanet \(\mathrm{GJ}\:436\mathrm{b}\), when we match an assumed tidal Love number k 2 value.  相似文献   

2.
Small tidal forces in the Earth–Moon system cause detectable changes in the orbit. Tidal energy dissipation causes secular rates in the lunar mean motion n, semimajor axis a, and eccentricity e. Terrestrial dissipation causes most of the tidal change in n and a, but lunar dissipation decreases eccentricity rate. Terrestrial tidal dissipation also slows the rotation of the Earth and increases obliquity. A tidal acceleration model is used for integration of the lunar orbit. Analysis of lunar laser ranging (LLR) data provides two or three terrestrial and two lunar dissipation parameters. Additional parameters come from geophysical knowledge of terrestrial tides. When those parameters are converted to secular rates for orbit elements, one obtains dn/dt = \(-25.97\pm 0.05 ''/\)cent\(^{2}\), da/dt = 38.30 ± 0.08 mm/year, and di/dt = ?0.5 ± 0.1 \(\upmu \)as/year. Solving for two terrestrial time delays and an extra de/dt from unspecified causes gives \(\sim \) \(3\times 10^{-12}\)/year for the latter; solving for three LLR tidal time delays without the extra de/dt gives a larger phase lag of the N2 tide so that total de/dt = \((1.50 \pm 0.10)\times 10^{-11}\)/year. For total dn/dt, there is \(\le \)1 % difference between geophysical models of average tidal dissipation in oceans and solid Earth and LLR results, and most of that difference comes from diurnal tides. The geophysical model predicts that tidal deceleration of Earth rotation is \(-1316 ''\)/cent\(^{2}\) or 87.5 s/cent\(^{2}\) for UT1-AT, a 2.395 ms/cent increase in the length of day, and an obliquity rate of 9 \(\upmu \)as/year. For evolution during past times of slow recession, the eccentricity rate can be negative.  相似文献   

3.
A family of well behaved perfect fluid balls has been derived starting with the metric potential g 44=B(1+Cr 2) n for all positive integral values of n. For n≥4, the members of this family are seen to satisfy the various physical conditions e.g. c 2 ρ≥p≥0,dp/dr<0,dρ/dr<0, along with the velocity of sound \((\sqrt{dp/c^{2}d\rho} )< 1\) and the adiabatic index ((p+c 2 ρ)/p)(dp/(c 2 dρ))>1. Also the pressure, energy density, velocity of sound and ratio of pressure and energy density are of monotonically decreasing towards the pressure free interface (r=a). The fluid balls join smoothly with the Schwarzschild exterior model at r=a. The well behaved perfect fluid balls so obtained are utilised to construct the superdense star models with their surface density 2×1014  gm/cm3. We have found that the maximum mass of the fluid balls corresponding to various values of n are decreasing with the increasing values of n. Over all maximum mass for the whole family turns out to be 4.1848M Θ and the corresponding radius as 19.4144 km while the red shift at the centre and red shift at surface as Z 0=1.6459 and Z a =0.6538 respectively this all happens for n=4. It is interesting to note that for higher values of n viz n≥170, the physical data start merging with that of Kuchowicz superdense star models and hence the family of fluid models tends to the Kuchowicz fluid models as n→∞. Consequently the maximum mass of the family of solution can not be less than 1.6096 M Θ which is the maximum mass occupied by the Kuchowicz superdense ball. Hence each member of the family for n≥4 provides the astrophysical objects like White dwarfs, Quark star, typical neutron star.  相似文献   

4.
In this paper, we consider the inverse problem of central configurations of n-body problem. For a given \({q=(q_1, q_2, \ldots, q_n)\in ({\bf R}^d)^n}\), let S(q) be the admissible set of masses denoted \({ S(q)=\{ m=(m_1,m_2, \ldots, m_n)| m_i \in {\bf R}^+, q}\) is a central configuration for m}. For a given \({m\in S(q)}\), let S m (q) be the permutational admissible set about m = (m 1, m 2, . . . , m n ) denoted
$S_m(q)=\{m^\prime | m^\prime\in S(q),m^\prime \not=m \, {\rm and} \, m^\prime\,{\rm is\, a\, permutation\, of }\, m \}.$
The main discovery in this paper is the existence of a singular curve \({\bar{\Gamma}_{31}}\) on which S m (q) is a nonempty set for some m in the collinear four-body problem. \({\bar{\Gamma}_{31}}\) is explicitly constructed by a polynomial in two variables. We proved:
  1. (1)
    If \({m\in S(q)}\), then either # S m (q) = 0 or # S m (q) = 1.
     
  2. (2)
    #S m (q) = 1 only in the following cases:
    1. (i)
      If s = t, then S m (q) = {(m 4, m 3, m 2, m 1)}.
       
    2. (ii)
      If \({(s,t)\in \bar{\Gamma}_{31}\setminus \{(\bar{s},\bar{s})\}}\), then either S m (q) = {(m 2, m 4, m 1, m 3)} or S m (q) = {(m 3, m 1, m 4, m 2)}.
       
     
  相似文献   

5.
New optimal, contractivity-preserving (CP), explicit, d-derivative, k-step Hermite–Obrechkoff series methods of order p up to \(p=20\), denoted by CP HO(d, k, p), with nonnegative coefficients are constructed. These methods are used to solve nonstiff first-order initial value problems \(y'=f(t,y)\), \(y(t_0)=y_0\). The upper bound \(p_u\) of order p of HO(d, k, p) can reach, approximately, as high as 2.4 times the number of derivatives d. The stability regions of HO(d, k, p) have generally a good shape and grow with decreasing \(p-d\). We, first, note that three selected CP HO methods: 4-derivative 7-step HO of order 13, denoted by HO(4, 7, 13), 5-derivative 6-step HO of order 13, denoted by HO(5, 6, 13), and 9-derivative 2-step HO of order 13, denoted by CMDAHO(13) compare favorably with Adams–Cowell of order 13, denoted by AC(13), in solving standard N-body problems over an interval of 1000 periods on the basis of the relative error of energy as a function of the CPU time. Next, the three HO methods compare positively with AC(13) in solving standard N-body problems on the basis of the growth of relative positional error and relative energy error over 10, 000 periods of integration. Finally, these three methods compare also well with P-stable methods of Cash and Franco et al. on some quasi periodic, second-order linear and nonlinear problems. The coefficients of selected HO methods are listed in the appendix.  相似文献   

6.
We present our B, V, Rc, and Ic observations of a \(3'.6 \times 3'\) field centered on the host galaxy of GRB 000926 (α2000.0=17h04m11s, \(\delta _{2000.0} = + 51^ \circ 47'9\mathop .\limits^{''} 8\)). The observations were carried out on the 6-m Special Astrophysical Observatory telescope using the SCORPIO instrument. The catalog of galaxies detected in this field includes 264 objects for which the signal-to-noise ratio is larger than 5 in each photometric band. The following limiting magnitudes in the catalog correspond to this limitation: 26.6 (B), 25.7 (V), 25.8 (R), and 24.5 (I). The differential galaxy counts are in good agreement with previously published CCD observations of deep fields. We estimated the photometric redshifts for all of the cataloged objects and studied the color variations of the galaxies with z. For luminous spiral galaxies with M(B)z～1.  相似文献   

7.
Zhanle Du 《Solar physics》2012,278(1):203-215
Smoothed monthly mean coronal mass ejection (CME) parameters (speed, acceleration, central position angle, angular width, mass, and kinetic energy) for Cycle 23 are cross-analyzed, showing that there is a high correlation between most of them. The CME acceleration (a) is highly correlated with the reciprocal of its mass (M), with a correlation coefficient r=0.899. The force (Ma) to drive a CME is found to be well anti-correlated with the sunspot number (R z), r=?0.750. The relationships between CME parameters and R z can be well described by an integral response model with a decay time scale of about 11 months. The correlation coefficients of CME parameters with the reconstructed series based on this model (\(\overline{r}_{\mathrm{f1}}=0.886\)) are higher than the linear correlation coefficients of the parameters with R z (\(\overline{r}_{\mathrm{0}}=0.830\)). If a double decay integral response model is used (with two decay time scales of about 6 and 60 months), the correlations between CME parameters and R z improve (\(\overline{r}_{\mathrm{f2}}=0.906\)). The time delays between CME parameters with respect to R z are also well predicted by this model (19/22=86%); the average time delays are 19 months for the reconstructed and 22 months for the original time series. The model implies that CMEs are related to the accumulation of solar magnetic energy. These relationships can help in understanding the mechanisms at work during the solar cycle.  相似文献   

8.
We construct a theory of the equilibrium figure and gravitational field of the Galilean satellite Io to within terms of the second order in the small parameter α. We show that to describe all effects of the second approximation, the equation for the figure of the satellite must contain not only the components of the second spherical function, but also the components of the third and fourth spherical functions. The contribution of the third spherical function is determined by the Love number of the third order h3, whose model value is 1.6582. Measurements of the third-order gravitational moments could reveal the extent to which the hydrostatic equilibrium conditions are satisfied for Io. These conditions are J3=C32=0 and C31/C33=?6. We have calculated the corrections of the second order of smallness to the gravitational moments J2 and C22. We conclude that when modeling the internal structure of Io, it is better to use the observed value of k2 than the moment of inertia derived from k2. The corrections to the lengths of the semiaxes of the equilibrium figure of Io are all positive and equal to ～64.5, ～26, and ～14 m for the a, b, and c axes, respectively. Our theory allows the parameters of the figure and the fourth-order gravitational moments that differ from zero to be calculated. For the homogeneous model, their values are:\(s_4 = \frac{{885}}{{224}}\alpha ^2 ,s_{42} = - \frac{{75}}{{224}}\alpha ^2 ,s_{44} = \frac{{15}}{{896}}\alpha ^2 ,J_4 = - \frac{{885}}{{224}}\alpha ^2 ,C_{42} = \frac{{75}}{{224}}\alpha ^2 ,C_{44} = \frac{{15}}{{896}}\alpha ^2 \).  相似文献   

9.
In this paper we study the relations between the main characteristics of groups and clusters of galaxies using the archival data of the SDSS and 2MASX catalogs. We have developed and implemented a new method of determining the size of galaxy systems and their effective radius which contains half of the galaxies and not half the luminosity, since the luminosity of the brightest galaxy in a group can account for over 50% of the total luminosity of the group. The derived parameters (log LK, logRe, and log σ200) for 94 systems of galaxies (0.0038 < z < 0.09) determine the Fundamental Plane (FP), which, with a scatter of 0.15, is similar in form to the FP of galaxy clusters obtained by Schaeffer et al. (1993) and D’Onofrio et al. (2013) with other methods and for different bands. We show that the FP in the near-infrared region (NIR) for 94 galaxy systems has the form of LK ∝ \(R_e^{0.70 \pm {{0.13}_\sigma }1.34 \pm 0.13}\), whereas in x-rays it has the form of—LX ∝ \(R_e^{1.15 \pm {{0.39}_\sigma }2.56 \pm 0.40}\). The form of the FP for groups and clusters is consistent with the FP for early-type galaxies determined in the same way. The form of the FP for galaxy systems deviates from the shape that one would expect from virial predictions. Adding the mass-to-light ratio as a fourth independent parameter has little effect on this deviation, but decreases the scatter of the FP for a sample of rich galaxy clusters by 12%.  相似文献   

10.
In this work, we first establish a simple procedure to obtain with 11-figure accuracy the values of Chandrasekhar’s H-function for isotropic scattering using a closed-form integral representation and the Gauss-Legendre quadrature. Based on the numerical values of the function produced by this method for various combinations of ? 0, the single scattering albedo, and μ, the cosine of the zenith angle θ of the direction of radiation emergent from or incident upon a semi-infinite scattering-absorbing medium, we propose a rational approximation formula with μ 1/4 and \(\sqrt{1-\varpi_{0}}\) as the independent variables. This allows us to reproduce the correct values of H(? 0,μ) within a relative error of 2.1×10?5 without recourse to any iterative procedure or root-finding process.  相似文献   

11.
  相似文献   

12.
We consider stars with radial velocities, proper motions, and distance estimates from the RAVE4 catalogue. Based on a sample of more than 145 000 stars at distances r < 0.5 kpc, we have found the following kinematic parameters: \({\left( {U,{\kern 1pt} V,{\kern 1pt} W} \right)_ \odot }\) = (9.12, 20.80, 7.66) ± (0.10, 0.10, 0.08) km s?1, Ω0 = 28.71 ± 0.63 km s?1 kpc?1, and Ω0 ’ = ?4.28 ± 0.11 km s?1 kpc?2. This gives the linear rotation velocity V 0 = 230 ± 12 km s?1 (for the adopted R 0 = 8.0 ± 0.4 kpc) and the Oort constants A = 17.12 ± 0.45 km s?1 kpc?1 and B = ?11.60 ± 0.77 km s?1 kpc?1. The 2D velocity distributions in the UV, UW, and VW planes have been constructed using a local sample, r < 0.25 kpc, consisting of ~47 000 stars. A difference of the UV velocity distribution from the previously known ones constructed from a smaller amount of data has been revealed. It lies in the fact that our distribution has an extremely enhanced branch near the Wolf 630 peak. A previously unknown peak at (U, V) = (?96, ?10) km s?1 and a separate new feature in the Wolf 630 stream, with the coordinates of its center being (U, V) = (30, ?40) km s?1, have been detected.  相似文献   

13.
We report the analysis of the young star clusters NGC 1960, NGC 2453 and NGC 2384 observed in the J (1.12 μm), H (1.65 μm) and K′ (2.2 μm) bands. Estimates of reddening, distance and age as E(B?V)=0.25, d=1380 pc and t=31.6 to 125 Myr for NGC 1960, E(B?V)=0.47, d=3311 pc and t=40 to 200 Myr for NGC 2453 and E(B?V)=0.25, d=3162 pc and t=55 to 125 Myr for NGC 2384 have been obtained. Also, we have extended the color–magnitude diagrams of these clusters to the fainter end and thus extended the luminosity functions to fainter magnitudes. The evolution of the main sequence and luminosity functions of these clusters have been compared with themselves as well as Lyngå 2 and NGC 1582.  相似文献   

14.
Spectrally resolved measurements of individual solar active regions (ARs) in the soft X-ray (SXR) range are important for studying dynamic processes in the solar corona and their associated effects on the Earth’s upper atmosphere. They are also a means of evaluating atomic data and elemental abundances used in physics-based solar spectral models. However, very few such measurements are available. We present spectral measurements of two individual ARs in the 0.5 to 2.5 nm range obtained on the NASA 36.290 sounding rocket flight of 21 October 2013 (at about 18:30 UT) using the Solar Aspect Monitor (SAM), a channel of the Extreme Ultaviolet Variability Experiment (EVE) payload designed for underflight calibrations of the orbital EVE on the Solar Dynamics Observatory (SDO). The EVE rocket instrument is a duplicate of the EVE on SDO, except the SAM channel on the rocket version was modified in 2012 to include a freestanding transmission grating to provide spectrally resolved images of the solar disk with the best signal to noise ratio for the brightest features, such as ARs. Calibrations of the EVE sounding rocket instrument at the National Institute of Standards and Technology Synchrotron Ultraviolet Radiation Facility (NIST/SURF) have provided a measurement of the SAM absolute spectral response function and a mapping of wavelength separation in the grating diffraction pattern. We discuss techniques (incorporating the NIST/SURF data) for determining SXR spectra from the dispersed AR images as well as the resulting spectra for NOAA ARs 11877 and 11875 observed on the 2013 rocket flight. In comparisons with physics-based spectral models using the CHIANTI v8 atomic database we find that both AR spectra are in good agreement with isothermal spectra (4 MK), as well as spectra based on an AR differential emission measure (DEM) included with the CHIANTI distribution, with the exception of the relative intensities of strong Fe?xvii lines associated with \(2p^{6}\)–\(2p^{5}3{s}\) and \(2p^{6}\)–\(2p^{5}3{d}\) transitions at about 1.7 nm and 1.5 nm, respectively. The ratio of the Fe?xvii lines suggests that the AR 11877 is hotter than the AR 11875. This result is confirmed with analysis of the active regions imaged by X-ray Telescope (XRT) onboard Hinode.  相似文献   

15.
We analyze the families of central configurations of the spatial 5-body problem with four masses equal to 1 when the fifth mass m varies from 0 to \(+\infty \). In particular we continue numerically, taking m as a parameter, the central configurations (which all are symmetric) of the restricted spatial (\(4+1\))-body problem with four equal masses and \(m=0\) to the spatial 5-body problem with equal masses (i.e. \(m=1\)), and viceversa we continue the symmetric central configurations of the spatial 5-body problem with five equal masses to the restricted (\(4+1\))-body problem with four equal masses. Additionally we continue numerically the symmetric central configurations of the spatial 5-body problem with four equal masses starting with \(m=1\) and ending in \(m=+\infty \), improving the results of Alvarez-Ramírez et al. (Discrete Contin Dyn Syst Ser S 1: 505–518, 2008). We find four bifurcation values of m where the number of central configuration changes. We note that the central configurations of all continued families varying m from 0 to \(+\infty \) are symmetric.  相似文献   

16.
An independent analysis of the molecular hydrogen absorption system at redshift z abs = 2.059 in the spectrum of the quasar J 2123?0050 is presented. The H2 system consists of two components (A and B) with column densities \(\log N_{{H_2}}^A = 17.94 \pm 0.01\) and \(N_{{H_2}}^B = 15.16 \pm 0.02\). The spectrum exhibits the lines of HDmolecules (logN HD A = 13.87±0.06) and the neutral speciesCI and Cl I associated with the H2 absorption system. For the molecular hydrogen lines near the quasar’s Lyβ and OVI emission lines, we detect a nonzero residual flux, ~3% of the total flux, caused by the effect of partial coverage of the quasar’s broad-line region by an H2 cloud. Due to the smallness of the residual flux, the effect does not affect the H2 column density being determined but increases the statistics of observations of the partial coverage effect to four cases. The uniqueness of the system being investigated is manifested in a high abundance of the neutral species H2 and CI at the lowest HI column density, logN HI = 19.18 ± 0.15, among the highredshift systems. The H2 and CI column densities in the system being investigated turn out to be higher than those in similar systems in our Galaxy and theMagellanic Clouds by two or three orders ofmagnitude. The \(N_{HD} /2N_{H_2 }\) ratio for component A has turned out to be also unusually high, (4.26 ± 0.60) × 10?5, which exceeds the deuterium abundance (D/H) for high-redshift systems by a factor of 1.5. Using the HI, H2, HD, and CI column densities as well as the populations of excited H2 and CI levels, we have investigated the physical conditions in components A and B. Component A represents the optically thick case; the gas has a low number density (~30 cm?3) and a temperature T ~ 140 K. In component B, the mediumis optically thin with n H ≤ 100 cm?3 and T ≥ 100 K. The ultraviolet (UV) background intensity in the clouds exceeds the mean intensity in our Galaxy by almost an order ofmagnitude. A high gas ionization fraction, \(n_{H^ + } /n_H \sim 10^{ - 2}\), which can be the result of partial shielding of the systemfrom hard UV radiation, is needed to describe the high HD and CI column densities. Using our simulations with the PDRMeudon code, we can reconstruct the observed column densities of the species within the model with a constant density (n H ~ 40 cm?3). A high H2 formation rate (higher than the mean Galactic value by a factor of 10?40) and high gas ionization fraction and UV background intensity are needed in this case.  相似文献   

17.
We present the results of the study of the eclipsing polar CRTS CSS081231 J071126+440405. Photometric observations allowed us to refine the orbital period of the system \(P_ \circ = 0_ \cdot ^d 0.08137673\). Considerable changes in the appearance of the object’s spectra have occurred over the period of September 20–21, 2001: the slope of the continuum changed from “red” to “blue”, and the variability of the line profiles over the duration of the orbital period has also changed. Doppler maps have shown a shift of the emission line-forming region along the accretion stream closer to the white dwarf. We measured the duration of the eclipse of the system and imposed constraints on the inclination angle \(78_ \cdot ^ \circ 7 < i < 79_ \cdot ^ \circ 3\). The derived radial velocity amplitude was used to obtain the basic parameters of the system: M1 = 0.86 ± 0.08M⊙, M2 = 0.18 ± 0.02 M⊙, q = 0.21 ± 0.01, RL2 = 0.20 ± 0.03 R⊙, A = 0.80 ± 0.03 R⊙. The spectra of the object exhibit cyclotron harmonics. Their comparison with model spectra allowed us to determine the parameters of the accretion column: B = 31–34 MG, Te = 10–12 keV, θ = 80–90°, and Λ = 105.  相似文献   

18.
The well-known Titius-Bode law (T-B) giving distances of planets from the Sun was improved by Basano and Hughes (1979) who found: $$a_n = 0.285 \times 1.523^n ;$$ a n being the semi-major axis expressed in astronomical units, of then-th planet. The integern is equal to 1 for Mercury, 2 for Venus etc. The new law (B-H) is more natural than the (T-B) one, because the valuen=?∞ for Mercury is avoided. Furthermore, it accounts for distances of all planets, including Neptune and Pluto. It is striking to note that this law:
  1. does not depend on physical parameters of planets (mass, density, temperature, spin, number of satellites and their nature etc.).
  2. shows integers suggesting an unknown, obscure wave process in the formation of the solar system.
In this paper, we try to find a formalism accounting for the B-H law. It is based on the turbulence, assumed to be responsible of accretion of matter within the primeval nebula. We consider the function $$\psi ^2 (r,t) = |u^2 (r,t) - u_0^2 |$$ , whereu 2(r, t) stands for the turbulence, i.e., the mean-square deviation velocities of particles at the pointr and the timet; andu 0 2 is the value of turbulence for which the accretion process of matter is optimum. It is obvious that Ψ2(r n,t0) = 0 forr n=0.285×1.523 n at the birth timet 0 of proto-planets. Under these conditions, it is easily found that $$\psi ^2 (r,t_0 ) = \frac{{A^2 }}{r}\sin ^2 [\alpha log r - \Phi (t_0 )]$$ With α=7.47 and Φ(t 0)=217.24 in the CGS system, the above function accounts for the B-H law. Another approach of the problem is made by considering fluctuations of the potentialU(r, t) and of the density of matter ρ(r, t). For very small fluctuations, it may be written down the Poisson equation $$\Delta \tilde U(r,t_0 ) + 4\pi G\tilde \rho (r,t_0 ) = 0$$ , withU(r, t)=U 0(r)+?(r, t 0 ) and \(\tilde \rho (r,t_0 )\) . It suffices to postulate \(\tilde \rho (r,t_0 ) = k[\tilde U(r,t_0 )/r^2 ](k = cte)\) for finding the solution $$\tilde U(r,t_0 ) = \frac{{cte}}{{r^{1/2} }}\cos [a\log r - \zeta (t_0 )]$$ . Fora=14.94 and ζ(t 0)=434.48 in CGS system, the successive maxima of ?(r,t 0) account again for the B-H law. In the last approach we try to write Ψ(r, t) under a wave function form $$\Psi ^2 (r,t) = \frac{{A^2 }}{r}\sin ^2 \left[ {\omega \log \left( {\frac{r}{v} - t} \right)} \right].$$ It is emphasized that all calculations are made under mathematical considerations.  相似文献   

19.
The most used method to calculate the coronal electron temperature [\(T_{\mathrm{e}} (r)\)] from a coronal density distribution [\(n_{\mathrm{e}} (r)\)] is the scale-height method (SHM). We introduce a novel method that is a generalization of a method introduced by Alfvén (Ark. Mat. Astron. Fys. 27, 1, 1941) to calculate \(T_{\mathrm{e}}(r)\) for a corona in hydrostatic equilibrium: the “HST” method. All of the methods discussed here require given electron-density distributions [\(n_{\mathrm{e}} (r)\)] which can be derived from white-light (WL) eclipse observations. The new “DYN” method determines the unique solution of \(T_{\mathrm{e}}(r)\) for which \(T_{\mathrm{e}}(r \rightarrow \infty) \rightarrow 0\) when the solar corona expands radially as realized in hydrodynamical solar-wind models. The applications of the SHM method and DYN method give comparable distributions for \(T_{\mathrm{e}}(r)\). Both have a maximum [\(T_{\max}\)] whose value ranges between 1?–?3 MK. However, the peak of temperature is located at a different altitude in both cases. Close to the Sun where the expansion velocity is subsonic (\(r < 1.3\,\mathrm{R}_{\odot}\)) the DYN method gives the same results as the HST method. The effects of the other free parameters on the DYN temperature distribution are presented in the last part of this study. Our DYN method is a new tool to evaluate the range of altitudes where the heating rate is maximum in the solar corona when the electron-density distribution is obtained from WL coronal observations.  相似文献   

20.
We present photoelectric and spectral observations of a hot candidate proto-planetary nebula—early B-type supergiant with emission lines in spectrum—IRAS 19336-0400. The light and color curves display fast irregular brightness variations with maximum amplitudes \(\Delta V = 0_ \cdot ^m 30\), \(\Delta B = 0_ \cdot ^m 35\), \(\Delta U = 0_ \cdot ^m 40\) and color-brightness correlations. By the variability characteristics IRAS 19336-0400 appears similar to other hot proto-planetary nebulae. Based on low-resolution spectra in the range λ4000–7500 Å we have derived absolute intensities of the emission lines Hα, Hβ, Hγ, [S II], [N II], physical conditions in gaseous nebula: n e = 104 cm?3, T e = 7000 ± 1000 K. The emission line Hα, Hβ equivalent widths are found to be considerably variable and related to light changes. By UBV-photometry and spectroscopy the color excess has been estimated: E B-V = 0.50–0.54. Joint photometric and spectral data analysis allows us to assume that the star variability is caused by stellar wind variations.  相似文献   

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