User:Tohline/Apps/OstrikerBodenheimerLyndenBell66

Rotationally Flattened White Dwarfs

Introduction

As we have reviewed in an accompanying discussion, Chandrasekhar (1935) was the first to construct models of spherically symmetric stars using the barotropic equation of state appropriate for a degenerate electron gas. In so doing, he demonstrated that the maximum mass of an isolated, nonrotating white dwarf is <math>M_3 = 1.44 (\mu_e/2)M_\odot</math>. A concise derivation of <math>~M_3</math> is presented in Chapter XI of Chandrasekhar (1967).

Something catastrophic should happen if mass is greater than <math>~M_3</math>. What will rotation do? Presumably it can increase the limiting mass.

• J. P. Ostriker, P. Bodenheimer & D. Lynden-Bell (1966), Phys. Rev. Letters, 17, 816: Equilibrium Models of Differentially Rotating Zero-Temperature Stars

… work by Roxburgh (1965, Z. Astrophys., 62, 134), Anand (1965, Proc. Natl. Acad. Sci. U.S., 54, 23), and James (1964, ApJ, 140, 552) shows that the [Chandrasekhar (1931, ApJ, 74, 81)] mass limit <math>~M_3</math> is increased by only a few percent when uniform rotation is included in the models, … In this Letter we demonstrate that white-dwarf models with masses considerably greater than <math>~M_3</math> are possible if differential rotation is allowed … models are based on the physical assumption of an axially symmetric, completely degenerate, self-gravitating fluid, in which the effects of viscosity, magnetic fields, meridional circulation, and relativistic terms in the hydrodynamical equations have been neglected.

Solution Strategy

When our objective is to construct steady-state equilibrium models of rotationally flattened, axisymmetric configurations, the accompanying introductory chapter shows how the overarching set of principal governing equations can be reduced in form to the following set of three coupled ODEs (expressed either in terms of cylindrical or spherical coordinates):

Cylindrical Coordinate Base

Spherical Coordinate Base

Poisson Equation <math>~ \frac{1}{\varpi} \frac{\partial }{\partial\varpi} \biggl[ \varpi \frac{\partial \Phi}{\partial\varpi} \biggr] + \frac{\partial^2 \Phi}{\partial z^2} </math> <math>~=</math> <math>~4\pi G \rho </math> The Two Relevant Components of the Euler Equation <math>~{\hat{e}}_\varpi</math>:     <math>~0</math> <math>~=</math> <math>~ \biggl[ \frac{1}{\rho}\frac{\partial P}{\partial\varpi} + \frac{\partial \Phi}{\partial\varpi}\biggr] - \frac{j^2}{\varpi^3} </math> <math>~{\hat{e}}_z</math>:     <math>~0</math> <math>~=</math> <math>~ \biggl[ \frac{1}{\rho}\frac{\partial P}{\partial z} + \frac{\partial \Phi}{\partial z} \biggr] </math>

Poisson Equation <math>~ \frac{1}{r^2} \frac{\partial }{\partial r} \biggl[ r^2 \frac{\partial \Phi }{\partial r} \biggr] + \frac{1}{r^2 \sin\theta} \frac{\partial }{\partial \theta}\biggl(\sin\theta ~ \frac{\partial \Phi}{\partial\theta}\biggr) </math> <math>~=</math> <math>~4\pi G\rho</math> The Two Relevant Components of the Euler Equation <math>~{\hat{e}}_r</math>:     <math> ~0 </math> = <math> \biggl[ \frac{1}{\rho} \frac{\partial P}{\partial r}+ \frac{\partial \Phi }{\partial r} \biggr] - \biggl[ \frac{j^2}{r^3 \sin^2\theta} \biggr] </math> <math>~{\hat{e}}_\theta</math>:     <math> ~0 </math> = <math> \biggl[ \frac{1}{\rho r} \frac{\partial P}{\partial\theta} + \frac{1}{r} \frac{\partial \Phi}{\partial\theta} \biggr] - \biggl[ \frac{j^2}{r^3 \sin^3\theta} \biggr] \cos\theta </math>

See Also

• Our discussion of Rotating White Dwarfs: Example Equilibria
• Fission Theory of Binary Star Formation

© 2014 - 2021 by Joel E. Tohline |   H_Book Home   |   YouTube   | Appendices: | Equations | Variables | References | Ramblings | Images | myphys.lsu | ADS | Recommended citation:   Tohline, Joel E. (2021), The Structure, Stability, & Dynamics of Self-Gravitating Fluids, a (MediaWiki-based) Vistrails.org publication, https://www.vistrails.org/index.php/User:Tohline/citation