In astrophysics, Chandrasekhar's white dwarf equation is an initial value ordinary differential equation introduced by the Indian American astrophysicist Subrahmanyan Chandrasekhar,[1] in his study of the gravitational potential of completely degenerate white dwarf stars. The equation reads as[2]
![{\displaystyle {\frac {1}{\eta ^{2}}}{\frac {d}{d\eta }}\left(\eta ^{2}{\frac {d\varphi }{d\eta }}\right)+(\varphi ^{2}-C)^{3/2}=0}](https://wikimedia.org/api/rest_v1/media/math/render/svg/083562d570529ff9d4907017cc95c812ec4b2625)
with initial conditions
![{\displaystyle \varphi (0)=1,\quad \varphi '(0)=0}](https://wikimedia.org/api/rest_v1/media/math/render/svg/4000bba512addfd0f9700d3a3d8617ba6efb5af8)
where
measures the density of white dwarf,
is the non-dimensional radial distance from the center and
is a constant which is related to the density of the white dwarf at the center. The boundary
of the equation is defined by the condition
![{\displaystyle \varphi (\eta _{\infty })={\sqrt {C}}.}](https://wikimedia.org/api/rest_v1/media/math/render/svg/3239adb8fcf5ff3fb7cdc36e9cf16a04d6c64d9d)
such that the range of
becomes
. This condition is equivalent to saying that the density vanishes at
.
Derivation
From the quantum statistics of a completely degenerate electron gas (all the lowest quantum states are occupied), the pressure and the density of a white dwarf are calculated in terms of the maximum electron momentum
standardized as
, with pressure
and density
, where
![{\displaystyle {\begin{aligned}&A={\frac {\pi m_{e}^{4}c^{5}}{3h^{3}}}=6.02\times 10^{21}{\text{ Pa}},\\&B={\frac {8\pi }{3}}m_{p}\mu _{e}\left({\frac {m_{e}c}{h}}\right)^{3}=9.82\times 10^{8}\mu _{e}{\text{ kg/m}}^{3},\\&f(x)=x(2x^{2}-3)(x^{2}+1)^{1/2}+3\sinh ^{-1}x,\end{aligned}}}](https://wikimedia.org/api/rest_v1/media/math/render/svg/6e8e4b0ce5617eb238ce23f6a13ed290b7eec1e8)
is the mean molecular weight of the gas, and
is the height of a small cube of gas with only two possible states.
When this is substituted into the hydrostatic equilibrium equation
![{\displaystyle {\frac {1}{r^{2}}}{\frac {d}{dr}}\left({\frac {r^{2}}{\rho }}{\frac {dP}{dr}}\right)=-4\pi G\rho }](https://wikimedia.org/api/rest_v1/media/math/render/svg/b2ae1df400f57e733bdaaf5a6b9e1792d9bbc268)
where
is the gravitational constant and
is the radial distance, we get
![{\displaystyle {\frac {1}{r^{2}}}{\frac {d}{dr}}\left(r^{2}{\frac {d{\sqrt {x^{2}+1}}}{dr}}\right)=-{\frac {\pi GB^{2}}{2A}}x^{3}}](https://wikimedia.org/api/rest_v1/media/math/render/svg/5b61a56b594046d76466f76736e6c937f8d7625b)
and letting
, we have
![{\displaystyle {\frac {1}{r^{2}}}{\frac {d}{dr}}\left(r^{2}{\frac {dy}{dr}}\right)=-{\frac {\pi GB^{2}}{2A}}(y^{2}-1)^{3/2}}](https://wikimedia.org/api/rest_v1/media/math/render/svg/da2f832ee90d7507fa757b1841aba75951da8da3)
If we denote the density at the origin as
, then a non-dimensional scale
![{\displaystyle r=\left({\frac {2A}{\pi GB^{2}}}\right)^{1/2}{\frac {\eta }{y_{o}}},\quad y=y_{o}\varphi }](https://wikimedia.org/api/rest_v1/media/math/render/svg/fdc366f12b4696f24b290465d1445252358c6b62)
gives
![{\displaystyle {\frac {1}{\eta ^{2}}}{\frac {d}{d\eta }}\left(\eta ^{2}{\frac {d\varphi }{d\eta }}\right)+(\varphi ^{2}-C)^{3/2}=0}](https://wikimedia.org/api/rest_v1/media/math/render/svg/083562d570529ff9d4907017cc95c812ec4b2625)
where
. In other words, once the above equation is solved the density is given by
![{\displaystyle \rho =By_{o}^{3}\left(\varphi ^{2}-{\frac {1}{y_{o}^{2}}}\right)^{3/2}.}](https://wikimedia.org/api/rest_v1/media/math/render/svg/4a22370477ba03d731c6ebc64ab99d3db62e3adb)
The mass interior to a specified point can then be calculated
![{\displaystyle M(\eta )=-{\frac {4\pi }{B^{2}}}\left({\frac {2A}{\pi G}}\right)^{3/2}\eta ^{2}{\frac {d\varphi }{d\eta }}.}](https://wikimedia.org/api/rest_v1/media/math/render/svg/b32a187efde33ebc948a605177f86ecd57b11101)
The radius-mass relation of the white dwarf is usually plotted in the plane
-
.
Solution near the origin
In the neighborhood of the origin,
, Chandrasekhar provided an asymptotic expansion as
![{\displaystyle {\begin{aligned}\varphi ={}&1-{\frac {q^{3}}{6}}\eta ^{2}+{\frac {q^{4}}{40}}\eta ^{4}-{\frac {q^{5}(5q^{2}+14)}{7!}}\eta ^{6}\\[6pt]&{}+{\frac {q^{6}(339q^{2}+280)}{3\times 9!}}\eta ^{8}-{\frac {q^{7}(1425q^{4}+11346q^{2}+4256)}{5\times 11!}}\eta ^{10}+\cdots \end{aligned}}}](https://wikimedia.org/api/rest_v1/media/math/render/svg/1d6fa7308fcd772f12cbdf2e61b47a44dcec18d7)
where
. He also provided numerical solutions for the range
.
Equation for small central densities
When the central density
is small, the equation can be reduced to a Lane–Emden equation by introducing
![{\displaystyle \xi ={\sqrt {2}}\eta ,\qquad \theta =\varphi ^{2}-C=\varphi ^{2}-1+x_{o}^{2}+O(x_{o}^{4})}](https://wikimedia.org/api/rest_v1/media/math/render/svg/7815f7fcd1bcc771b5acd0dda7d38455e0afc1ce)
to obtain at leading order, the following equation
![{\displaystyle {\frac {1}{\xi ^{2}}}{\frac {d}{d\xi }}\left(\xi ^{2}{\frac {d\theta }{d\xi }}\right)=-\theta ^{3/2}}](https://wikimedia.org/api/rest_v1/media/math/render/svg/be6384c24da90b75b830748f8928d9adb00b3986)
subjected to the conditions
and
. Note that although the equation reduces to the Lane–Emden equation with polytropic index
, the initial condition is not that of the Lane–Emden equation.
Limiting mass for large central densities
When the central density becomes large, i.e.,
or equivalently
, the governing equation reduces to
![{\displaystyle {\frac {1}{\eta ^{2}}}{\frac {d}{d\eta }}\left(\eta ^{2}{\frac {d\varphi }{d\eta }}\right)=-\varphi ^{3}}](https://wikimedia.org/api/rest_v1/media/math/render/svg/f2799b12b54a8739d6359f12189abda3d048c4d0)
subjected to the conditions
and
. This is exactly the Lane–Emden equation with polytropic index
. Note that in this limit of large densities, the radius
![{\displaystyle r=\left({\frac {2A}{\pi GB^{2}}}\right)^{1/2}{\frac {\eta }{y_{o}}}}](https://wikimedia.org/api/rest_v1/media/math/render/svg/d24a7ddb64a598111942ab5c1be143f28250c20a)
tends to zero. The mass of the white dwarf however tends to a finite limit
![{\displaystyle M\rightarrow -{\frac {4\pi }{B^{2}}}\left({\frac {2A}{\pi G}}\right)^{3/2}\left(\eta ^{2}{\frac {d\varphi }{d\eta }}\right)_{\eta =\eta _{\infty }}=5.75\mu _{e}^{-2}M_{\odot }.}](https://wikimedia.org/api/rest_v1/media/math/render/svg/50e01ac20edfa2adf8df13d09451ae4a8bd80af0)
The Chandrasekhar limit follows from this limit.
See also
References
- ^ Chandrasekhar, Subrahmanyan, and Subrahmanyan Chandrasekhar. An introduction to the study of stellar structure. Vol. 2. Chapter 11 Courier Corporation, 1958.
- ^ Davis, Harold Thayer (1962). Introduction to Nonlinear Differential and Integral Equations. Courier Corporation. ISBN 978-0-486-60971-3.