URCA shells in dense stellar interiors

Thermal and vibrational energy losses due to URCA shells in stellar interiors are calculated. Analytic expressions are derived for semidegenerate, relativistic electrons. Results are given for more general cases calculated with a computer. The calculations are carried out for a large number of nuclei that may contribute to URCA energy losses in various stages of stellar evolution. An illustration is given of the cooling and vibrational damping of a white dwarf. For a central Fermi energy 5 MeV, the internal temperature of the star should be reduced to the order of 10⁶ K and the relative vibrational amplitude should be reduced to the order of 10^–5 on a time scale of 10⁹ yr. URCA shells are present, the URCA neutrino energy loss dominates in the temperature region up to about 2×10⁹ K.     

The ordinary URCA process in the presence of positrons and large magnetic fields

The effect of positron capture on the ordinary URCA neutrino luminosity in a zero magnetic field is investigated for several values of the degeneracy parameter and the range of temperatures 5×10⁸K–5×10¹⁰K. The rate for this process is then compared with those in large magnetic fields (on the order ofH _c =m ² c ³/eh=4.414×10¹⁰ G). The results indicate that positron capture reduces the effect of large magnetic fields on this process at high temperatures.     

Neutrino luminosity by the ordinary URCA process in an intense magnetic field

The neutrino luminosity by the ordinary URCA process in a strongly magnetized electron gas is computed. General formulae are presented for the URCA energy loss rates for an arbitrary degree of degeneracy. Analytic expressions are derived for a completely degenerate, relativistic electron plasma in the special case of neutron-proton conversion. Numerical results are given for more general cases.The main results are as follows: the URCA energy loss rates are drastically reduced for the regime of great degeneracy by a factor up to 10^–3 for 1, andT ₉10, where =H/H _q ,H _q =m ² c ³/eh=4.414×10¹³ G. In the non-degenerate regime the neutrino luminosity is enhanced approximately linearly with for the temperature range 1T ₉10. Possible applications to white dwarfs and neutron stars are briefly discussed.We have been recently informed that in Gamow home-dialect (Odessa dialect) URCA means thief — (Private communication from Prof. G. Wataghin).     

The heating of the solar plasma due to microwave phenomena correlated with type II meter bursts

The heating of the solar plasma of those layers is considered where the microwave bursts are emitted. In a first step, we restrict ourselves to phenomena correlated with the so-called type II m bursts. Bursts of this kind are excited by shock-waves initiated near the optical flare region. These shock-waves spread out into the higher corona, and if the shock strength is sufficiently high, the microwave region is heated to 10⁷ K. But this temperature is too low to explain the burst radiation. In this paper, it is shown that at plasma temperatures about 10⁷ K a fairly high number of electrons is accelerated by Alfvén waves to equivalent kinetic temperatures of about 10⁸ K. We assume that the Alfvén waves are generated near the sunspots, and, therefore, the accelerated electrons run along the magnetic-field lines into the microwave source lying between the two spots of an assumed dipole field. Within this source, the considered electrons thermalize and, after a short time, the source reaches temperatures of 5 × 10⁷ K to 10⁸ K.A plasma of this temperature with an electron density about 5 × 10⁹ cm^–3 and a magnetic induction of 300 G is optically thick even at frequencies about 10 GHz, because the gyromagnetic absorption is very high.     

Thermal equilibria of coronal magnetic loops

Equations of thermal equilibrium along coronal loops with footpoint temperatures of 2 × 10⁴ K are solved. Three fundamentally different categories of solution are found, namely hot loops with summit temperatures above about 4 × 10⁵ K, cool loops which are cooler than 8 × 10⁴ K along their whole length and hot-cool loops which have summit temperatures around 2 × 10⁴ K but much hotter parts at intermediate points between the summit and the footpoints. Hot loops correspond to the hot corona of the Sun. The cool loops are of relevance for fibrils, for the cool cores observed by Foukal and also for active-region prominences where the magnetic field is directed mainly along the prominence. Quiescent prominences consist of many cool threads inclined to the prominence axis, and each thread may be modelled as a hot-cool loop. In addition, it is possible for warm loops at intermediate summit temperatures (8 × 10⁴K to 4 × 10⁵ K) to exist, but the observed differential emission measure suggests that most of the plasma in the solar atmosphere is in either the hot phase or the cool phase. Thermal catastrophe may occur when the length or pressure of a loop is so small that the hot solution ceases to exist and there are only cool loop solutions. Many loops can be superimposed to form a coronal arcade which contains loops of several different types.     

A numerical test of the quasi-equilibrium approximation for stellar silicon burning

A measure of the range of the validity of the nuclear quasi-equilibrium approximation, employed byBodansky, Clayton andFowler (1968) for the treatment of stellar silicon burning, is obtained by comparison of the predicted abundances with the results of a numerical solution of the equations for the time rates of change of the nuclear abundances. The results of these calculations, performed for a temperatureT=3×10⁹K, are compared with those obtained by Bodanskyet al. for a temperatureT=5×10⁹K. The time required for the realization of the quasi-equilibrium condition at high temperatures is found to comprise a more substantial fraction of the silicon burning lifetime. This behavior is found to be attributable to the relative temperature sensitivities of the nuclear photodisintegration rates which determine the silicon burning lifetime and the charged particle reaction rates (largely Ti⁴⁴(,p) V⁴⁷) which determine the rate of buildup of iron peak nuclei.     

Morphology and spatial distribution of XUV and X-ray emissions in an active region observed from Skylab

We studied the morphology and spatial distribution of loops in an active region, using coordinated observations obtained with both the S082A XUV spectroheliograph and the S056 grazingincidence X-ray telescope on Skylab. The active region loops in the temperature range 5 × 10⁵ –3 × 10⁶ K fall basically into two distinctive groups: the hot loops with temperatures 2–3 × 10⁶ K as observed in coronal lines and X-rays, and the relatively cool loops with temperature 5 × 10⁵ –1 × 10⁶ K as observed in transition-zone lines (Ne vii, Mg ix). The brightest hot coronal loops in the active region are mostly low-lying, compact, closely-packed, and show greater stability than the transition-zone loops, which are fewer in number, large, and slender. The observed aspect ratio of the hot coronal loops is in the range of 0.1 and 0.2, which are almost two orders of magnitude larger than those for the Ne vii loops. Brief discussion of the MHD stability of the loops in terms of the aspect ratio is presented.     

Rapid neutron capture process in supernovae and chemical element formation

The rapid neutron capture process (r-process) is one of the major nucleosynthesis processes responsible for the synthesis of heavy nuclei beyond iron. Isotopes beyond Fe are most exclusively formed in neutron capture processes and more heavier ones are produced by the r-process. Approximately half of the heavy elements with mass number A > 70 and all of the actinides in the solar system are believed to have been produced in the r-process. We have studied the r-process in supernovae for the production of heavy elements beyond A = 40 with the newest mass values available. The supernova envelopes at a temperature >109 K and neutron density of 10²⁴ cm^?3 are considered to be one of the most potential sites for the r-process. The primary goal of the r-process calculations is to fit the global abundance curve for solar system r-process isotopes by varying time dependent parameters such as temperature and neutron density. This method aims at comparing the calculated abundances of the stable isotopes with observation. We have studied the r-process path corresponding to temperatures ranging from 1.0 × 10⁹ K to 3.0 × 10⁹ K and neutron density ranging from 10²⁰ cm^?3 to 10³⁰ cm^?3. With temperature and density conditions of 3.0 × 10⁹ K and 10²⁰ cm^?3 a nucleus of mass 273 was theoretically found corresponding to atomic number 115. The elements obtained along the r-process path are compared with the observed data at all the above temperature and density range.     

Ultra-high energy neutrino fluxes from supermassive AGN black holes

We compute the ultra-high energy (UHE) neutrino fluxes from plausible accreting supermassive black holes closely linking to the 377 active galactic nuclei (AGNs). They have well-determined black hole masses collected from the literature. The neutrinos are produced via simple or modified URCA processes, even after the neutrino trapping, in superdense proto-matter medium. The resulting fluxes are ranging from: (1) (quark reactions)— $J^{q}_{\nu\varepsilon}/(\varepsilon_{d}\ \mathrm{erg}\,\mathrm{cm}^{-2}\,\mathrm{s}^{-1}\,\mathrm{sr}^{-1})\simeq8.29\times 10^{-16}$ to 3.18×10^?4, with the average $\overline{J}^{q}_{\nu\varepsilon}\simeq5.53\times 10^{-10}\varepsilon_{d}\ \mathrm{erg}\,\mathrm{cm}^{-2}\,\mathrm{s}^{-1}\,\mathrm{sr}^{-1}$ , where ε _d ～10^?12 is the opening parameter; (2) (pionic reactions)— $J^{\pi}_{\nu\varepsilon} \simeq0.112J^{q}_{\nu\varepsilon}$ , with the average $J^{\pi}_{\nu\varepsilon} \simeq3.66\times 10^{-11}\varepsilon_{d}\ \mathrm{erg}\,\mathrm{cm}^{-2}\,\mathrm{s}^{-1}\,\mathrm{sr}^{-1}$ ; and (3) (modified URCA processes)— $J^{URCA}_{\nu\varepsilon}\simeq7.39\times10^{-11} J^{q}_{\nu\varepsilon}$ , with the average $\overline{J}^{URCA}_{\nu\varepsilon} \simeq2.41\times10^{-20} \varepsilon_{d}\ \mathrm{erg}\,\mathrm{cm}^{-2}\,\mathrm{s}^{-1}\,\mathrm{sr}^{-1}$ . We conclude that the AGNs are favored as promising pure neutrino sources, because the computed neutrino fluxes are highly beamed along the plane of accretion disk, peaked at high energies and collimated in smaller opening angle θ～ε _d .     

Electron densities derived from line intensity ratios: Beryllium isoelectronic sequence

A direct method for determining electron densities from emission line intensities of ions in the beryllium isoelectronic sequence is described and then applied to the analysis of extreme ultraviolet Ciii and Ov spectra from both quiet and active areas in the solar transition region. The results are consistent with a value of N _e T _e = 6 × 10¹⁴ cm^-3K for the quiet Sun at temperatures of 5 × 10⁴ to 3 × 10⁵K. Electron densities are approximately five times greater in active regions than in the quiet Sun.     

The relation between hard X-ray and transition-region line emission in solar flares

Observational evidence suggests that both the hard X-ray and ultraviolet emission from the impulsive phase of flares result from an electron beam. We present the results of model calculations that are consistent with this theory. The impulsive phase is envisioned as occurring in many small magnetically confined loops, each of which maintains an electron beam for only a few seconds. This model successfully matches several observed aspects of the impulsive phase. The corona is heated to less than 2 × 10⁶ K, maximum enhanced emission occurs in lines formed near 10⁵ K, and there is only slight enhancement between 10⁵ and 2 × 10⁶ K. The slope of the observed relationship between hard X-ray and Ov 1371 Å emission is also matched, but the relative emission is not. The calculations indicate that UV emission lines formed below a temperature of about 10⁵ K will arise predominantly from the chromospheric region heated by the electron beam to transition region temperatures. Emission lines formed at higher temperatures will be produced in the transition region. This should be detectable in density-sensitive line ratios. To account successfully for the impulsive UV emission, the peak temperature in the impulsively heated loops must remain below about 2 × 10⁶ K. Thus our model implies that the impulsive heating takes place in different loops from the hotter gradual phase emission.     

Prominence–corona transition region plasma diagnostics from SOHO observations

New results concerning prominence observations and in particular the prominence–corona transition region (PCTR) are presented. In order to cover a temperature range from 2 × 10⁴ to 7 × 10⁵ K, several emission lines in many different ionization states were observed with SUMER and CDS on board SOHO. EM and DEM were measured through the whole PCTR. We compared the prominence DEM with the DEM from other solar structures (active region, coronal hole and the chromosphere–corona transition region (CCTR)). We notice a displacement of the prominence DEM minimum towards lower temperatures with respect to the minimum of the other structures. Electron density and pressure diagnostics have been made from the observed C III lines. Local electron density and pressure for T ∼ 7 × 10⁴ K are respectively log N _e = 9.30_−0.34 ^+0.30 and 0.0405_−0.014 ^+0.012. Extrapolations over the entire PCTR temperature range are in good agreement with previous SOHO results (Madjarska et al., 1999). We also provide values of electron density and pressure in two different regions of the prominence (center and edge). The Doppler velocity in the PCTR shows a trend to increase with temperature (at least up to 30 km s -1 at T ∼ 7 × 10⁴ K), an indication of important mass flows. A simple morphological model is proposed from density and motion diagnostics. If the prominence is taken as a magnetic flux tube, one can derive an opening of the field lines with increasing temperature. If the prominence is represented as a collection of threads, their number increases with temperature from 20 to 800. Derived filling factors can reach values as low as 10⁻³ for a layer thickness of the order of 5000 km. The variation of non-thermal velocities is determined for the first time, in the temperature range from 2 × 10⁴ to 7 × 10⁵ K. The quite clear similarity with the CCTR non-thermal velocities would indicate that heating mechanisms in the PCTR could be the same as in the CCTR (wave propagation, turbulence MHD).     

Formation of stellar associations

The Rayleigh-Taylor instability forms massive complexes. When 10²¹ atoms cm^–2 are gathered, X-rays which heat the gas and UV-rays which ionize carbon are absorbed. A layer should appear with temperatures as low as 6 K and density to 4×10³ cm^–3. Finally the layer is fragmented into stars whose masses may even be less than one solar mass. The temperature of the layer should increase with time because part of free carbon is gradually absorbed by dust. Therefore more massive stars should appear after less massive stars. The stars which are formed kept near the layer by its gravitation. When their total mass becomes comparable with the mass of the layer, they should fall to the galactic plane in agreement with observed proper motions of several studied stellar systems.     

Isotopic r-process abundances produced by supernova explosions

The rapid neutron capture process (r-process) is one of the major nucleosynthesis processes responsible for the synthesis of heavy nuclei beyond iron. Isotopes beyond Fe are most exclusively formed in neutron capture processes and more heavier ones are produced by the r-process. Approximately half of the heavy elements with mass number A>70 and all of the actinides in the solar system are believed to have been produced in the r-process. We have studied the r-process in supernovae for production of heavy elements beyond A=40 with the newest mass values available. The supernovae envelopes at a temperature >10⁹ K and neutron density of 10²⁴ cm⁻³ are considered to be one of the most potential sites for the r-process. We investigate the r-process in a site-independent, classical approach which assumes a chemical equilibrium between neutron captures and photodisintegrations followed by a β-flow equilibrium. We have studied the r-process path corresponding to temperatures ranging from 1.0×10⁹ K to 3.0×10⁹ K and neutron density ranging from 10²⁰ cm⁻³ to 10³⁰ cm⁻³. The primary goal of the r-process calculations is to fit the global abundance curve for solar system r-process isotopes by varying time dependent parameters such as temperature and neutron density. This method aims at comparing the calculated abundances of the stable isotopes with observation. The abundances obtained are compared with supernova explosion condition and found in good agreement. The elements obtained along the r-process path are compared with the observed data at all the above temperature and density range.     

Explosive hydrogen burning in supernovae

Rapid proton capture is supposed to be responsible for the synthesis of a number of proton-rich nuclei. This process of hydrogen burning is considered here for mass elements, the atomic numbers of which range fromZ=10 toZ=20. The possible site for this process is assumed to be the outer envelope of the supernova at a proton number density (n _p )ranging fromn _p =10²² cm^–3 ton _p =10²⁸ cm^–3 at temperatures in the range ofT=2–3×10⁹ K.The capture path is determined by considering that a dynamical equilibrium between (p, ) and (,p) reactions exists between the reacting nuclei. In this situation, the abundances of elements become proportional to the lifetime of ⁺ decaying nuclei at the waiting points.It is suggested that these rapid proton-capture reactions are responsible for the production of a number of nuclei in the rangeA40 during supernova outbursts.     

Steady mass-loss from supermassive stars

Structures of Newtonian super-massive stars are calculated with the opacity for Comptor effectK ₀/(1 + T), whereK ₀=0.21(1 +X and =2.2×10^–9K^–1. The track of the Main-Sequence is turned right in the upper part of the HR diagram. Mass loss will occur in a Main-Sequence stage for a star with mass larger than a critical mass. The cause of mass loss and the expansion of the radius is continuum radiation pressure. The critical mass for mass loss is 1.02×10⁶ M for a Population I star, and 1.23×10⁵ M for Population III star. Mass loss rates expected in these stars are 3.3×10^–3 and 4.0×10^–3 M yr^–1, respectively.Paper presented at the IAU Third Asian-Pacific Regional Meeting, held in Kyoto, Japan, between 30 September–6 October, 1984.     

Mass loss from solar-type stars

Winds are directly detected from solar-type stars only when they are very young. At ages 10⁶ yr, these stars have mass loss rates 10⁶ times the mass flux of the present solar wind. Although these young T Tauri stars exhibit ultraviolet transition-region and X-ray coronal emission, the large particle densities of the massive winds lead to efficient radiative cooling, and wind temperatures are only 10⁴ K. In these circumstances thermal acceleration is unlikely to play an important role in driving the mass loss. Turbulent energy fluxes may be responsible for the observed mass loss, particularly if substantial magnetic fields are present.The presence of stellar mass loss is indirectly shown by the spindown of low-mass stars as they age. It appears that many solar-mass stars spin up as they contract toward the Main-Sequence, reaching a maximum equatorial velocity of 50 to 100 km s^–1. These stars spin down rapidly upon reaching the Main Sequence. Spindown may be enhanced by a decoupling or lag between convective envelope and radiative core. Because this spindown occurs fairly early in a solar-type star's history, the internal structure of old stars like the Sun may not depend upon initial conditions.     

The effect of UV stars on the intergalactic medium

We have investigated the effect of ionizing radiation from the UV stars (hot prewhite dwarfs) on the intergalactic medium (IGM). If the UV stars are powered only by gravitational contraction they radiate most of their energy at a typical surface temperature of 1.5×10⁵ K which produces a very highly ionized IGM in which the elements carbon, nitrogen and oxygen are left with only one or two electrons. This results in these elements being very inefficient coolants. The gas is cooled principally by free-free emission and the collisional ionization of hydrogen and helium. For a typical UV star temperature ofT=1.5×10⁵ K, the temperature of the ionized gas in the IGM isT _g =1.2×10⁵ K for a Hubble constantH _o=75 km s^–1 Mpc^–1 and a hydrogen densityn _H =10^–6 cm^–3. Heating by cosmic rays and X-rays is insignificant in the IGM except perhaps inHi clouds because when a hydrogen atom recombines in the IGM it is far more likely to be re-ionized by a UV-star photon than by of the other two types of particles due to the greater space density of UV-star photons and their appreciably larger ionization cross-sections. If the UV stars radiate a substantial fraction of their energy in a helium-burning stage in which they have surface temperatures of about 5×10⁴ K, the temperature of the IGM could be lowered to about 5×10⁴ K.     

Soft X-ray enhancement during flares

The estimates of quiescent and flare time temperatures of soft X-ray emitting regions on the Sun are obtained for flares observed during March–August 1967 from X-ray observations in two soft X-ray bands, 2–12 Å (Explorer-33 data) and 8–12 Å (OSO-3 data). It is concluded that hot coronal condensation, originally at 2–3 × 10⁶ K, is raised to the temperature of about 4–5 × 10⁶ K and is responsible for soft X-ray enhancement.On leave from Physics Department, College of Engineering, Aurangabad, India.