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Stellar evolution:
As it collapses, a molecular cloud breaks into smaller and smaller pieces. In each of these fragments, the collapsing gas releases gravitational potential energy as heat. As its temperature and pressure increase, the fragments condense into rotating spheres of gas. Once the gas is hot enough for the internal pressure to support the fragment against further gravitational collapse (hydrostatic equilibrium), the object is known as a protostar [Aud72].
Accretion of material onto the protostar continues partially through a circumstellar disc.
When the density and temperature are high enough, deuterium fusion begins, and the outward pressure of the resultant radiation slows (but does not stop) the collapse. Material comprising the cloud continues to « rain » onto the protostar.
The protostar follows a Hayashi track on the Hertzsprung-Russell (H-R) diagram. The contraction will proceed until the Hayashi limit is reached, and thereafter contraction will continue with the temperature remaining stable. Stars with less than 0.5 solar masses thereafter join the main sequence. For more massive protostars, at the end of the Hayashi track they will slowly collapse in near hydrostatic equilibrium, Finally, hydrogen begins to fuse in the core of the star, and the rest of the enveloping material is cleared away. This ends the protostellar phase and begins the star’s main sequence phase on the H-R diagram. The later evolution of stars are studied in stellar evolution [Aud72]..
As the hydrogen around the core is consumed, the core absorbs the resulting helium, causing it to contract further, which in turn causes the remaining hydrogen to fuse even faster. This eventually leads to ignition of helium fusion (which includes the triple-alpha process) in the core.
In stars of more than approximately 0.5 solar masses, electron degeneracy pressure may delay helium fusion for millions or tens of millions of years; in more massive stars, the combined weight of the helium core and the overlying layers means that such pressure is not sufficient to delay the process significantly.
ASTROPHYSICAL NUCLEAR REACTIONS
The thermonuclear reactions of astrophysical interest concern mainly the capture of nucleons or alpha particles. A limited number of fusion reactions involving heavy ions (12C, 16O) are also of great importance. Charged-particle induced reactions are essential for the energy budget of a star, as well as for the production of new nuclides in stellar and non-stellar (Big Bang) situations. In contrast, the role of neutron captures is largely restricted to nucleosynthesis, their energetic impact being negligible.
The heart of stellar evolution and nucleosynthesis is the thermonuclear reactions. It is the fusion of light nuclei into heavier nuclei that liberates kinetic energy (at the expense of mass) and serves as the interior source of the energy radiated from the surface. The condition that the power liberated internally balance the power radiated from the surface determines a steady state in the structure of the star. That situation cannot be a truly static one, however, because the very reactions that liberate energy necessarily change the chemical composition of the stellar interior. It is the slow change of chemical composition that causes the structure of the star to the chemical composition of the interstellar medium will have been altered by the thermonuclear debris. Stated most simply, it is the working hypothesis of the stellar nucleosynthesist that all or part of the heavy elements found in our galaxy have been synthesized in the interiors of stars by these same fusion reactions.
A complete science of thermonuclear reaction rates is formidable. It involves complicated details of nuclear physics, many of which are still unsolved. The mechanism of each reaction must be scrutinized to achieve assurance of the proper prescription for the stellar reaction rate. Still there are a few basic physical principles that are common to the computation of all thermonuclear reaction rates.
INTRODUCTION
CHAPTER 1: ELEMENTS OF NUCLEAR ASTROPHYSICS
1.1.Introduction
1.1. The Hurtzprung‐ Russell Diagramm (HRD) Nucleosynthesis
1.2.1.Hydrogen burning
1.2.1.1. The proton‐proton chain
1.2.1.2. The CNO cycle
1.2.2. Heli um burning
1.2.3. The Carbon burning
1.2.4. The Oxygen burning
1.2 .5. Photodesintegration
1.2.6. r ,sand p process
1.2.7. Spallation re action
1.3. Stellar evolution
CHAPTER 2: ASTROPHYSICAL NUCLEAR REACTIONS
2.1. Introduction
2.2.Kinematics and energetics
2.3. Reac tion rate
2.4. The cross section
2.5. The astrophysical factor
CHAPTER 3: ATOMIC SCREENING IN ASTROPHYSICAL NUCLEAR REACTIONS
3.1. Electron screening effects in low energy reactions
3.1.1. Introduction
3 .2. Coulomb screened potential
3.2.1. Prior work
3.22. BES model (Bat na Electron Screening Model)
3.2.3. Comparison with Liolios Model
3.3 . One electron screening effects
3.4. Two electron screening effects
3.5. Screening effect with excited electrons
3.5.1. The projectile in the 2s state
3.5.2. Both atoms are in the 2s state
3.6. Astrophysical S (E) of at solar energy
3.6.1.Mechanism of interaction
3.6.2. Calculation of the screened astrophysical factor
3.6.2.1. The linear approximati on
3.6.2.2. The quadratic approximation
CHAPTER 4: MAGNETICALLY CATAL YZED SCREENING IN FUSION REACTIONS
4.1. Strong magnetic field in astro nomy
4.2. Motion of particle in a uniform magnetic field
4.2.1.Soluti on in the Cartesian coordinates
4.2.2. Solution in the cylindrical coordinates
4.3. Binding energy of Hydrogen atoms in strong magnetic field
4.3.1. Introduction
4.3.2. Hydrogen atom in strong magnetic field
a. Axial wave function
4.4. Screening potential
4.4.1. Introduction
4.4.2. Heyl’s potential
4.4.3. BES’s Potential
4.5. The acceleration factor
4.6. Heyl versus BES calculation
CHAPTER 5: SCREENED ALPHA DECAY IN SUPERSTRONG MAGNETIC FIELDS AND DENSE
ASTROPHYSICAL PLASMAS
5.1. Introduction
Basic alpha decay processes
5.2. Theory of a lpha emission
5.3. Screened alpha decay in terrestrial environment
5. 4. Magnetically catalyzed alpha decay in magnetars
5.4.1. The BMSC Poten tial
5.4.2. The magnetically enhanced screening factor
5.5. Screened alpha decay in dense astrophysical plasmas
5.5.1. The linear plasma shie lding
5.5.2. linear plasma schielding model
5.5.3. The non linear plasma screening
5.5.4. Effect of nonlinear plasma screening
Conclusion
References
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