# User:Tohline/SSC/Structure/LimitingMasses

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 Revision as of 07:15, 31 May 2013 (view source)Tohline (Talk | contribs) (→Mass Upper Limits)← Older edit Revision as of 18:57, 31 May 2013 (view source)Tohline (Talk | contribs) (→Mass Upper Limits: add to and modify discussion)Newer edit → Line 3: Line 3: {{LSU_HBook_header}} {{LSU_HBook_header}} + ==Preamble== Spherically symmetric, self-gravitating, equilibrium configurations can be constructed from gases exhibiting a wide variety of degrees of compressibility.  When examining how the internal structure of such configurations varies with compressibility, or when examining the relative stability of such structures, it can be instructive to construct models using a polytropic equation of state because the degree of compressibility can be adjusted by simply changing the value of the polytropic index, {{User:Tohline/Math/MP_PolytropicIndex}}, across the range, $0 \le n \le \infty$.  (Alternatively, one can vary the effective adiabatic exponent of the gas, $\gamma_g = 1 + 1/n$.)  In particular, $n = 0$ ($\gamma_g = \infty$) represents a ''hard'' equation of state and describes an incompressible configuration, while $n = \infty$ ($\gamma_g = 1$) represents an isothermal and extremely ''soft'' equation of state. Spherically symmetric, self-gravitating, equilibrium configurations can be constructed from gases exhibiting a wide variety of degrees of compressibility.  When examining how the internal structure of such configurations varies with compressibility, or when examining the relative stability of such structures, it can be instructive to construct models using a polytropic equation of state because the degree of compressibility can be adjusted by simply changing the value of the polytropic index, {{User:Tohline/Math/MP_PolytropicIndex}}, across the range, $0 \le n \le \infty$.  (Alternatively, one can vary the effective adiabatic exponent of the gas, $\gamma_g = 1 + 1/n$.)  In particular, $n = 0$ ($\gamma_g = \infty$) represents a ''hard'' equation of state and describes an incompressible configuration, while $n = \infty$ ($\gamma_g = 1$) represents an isothermal and extremely ''soft'' equation of state. - '''ISOLATED POLYTROPES:'''  [[User:Tohline/SSC/Structure/Polytropes#Polytropic_Spheres|Isolated polytropic spheres [*] ]] exhibit three attributes that are especially key in the context of our present discussion: + ==Isolated Polytropes== + [[User:Tohline/SSC/Structure/Polytropes#Polytropic_Spheres|Isolated polytropic spheres]] exhibit three attributes that are especially key in the context of our present discussion:
1. The equilibrium structure is dynamically stable if $n < 3$.
2. The equilibrium structure is dynamically stable if $n < 3$. Line 14: Line 16:
- '''ISOTHERMAL SPHERES:'''  Isothermal spheres (polytropes with index $n=\infty$) are discussed in a wide variety of astrophysical contexts because it is not uncommon for physical conditions to conspire to create an extended volume throughout which a configuration exhibits uniform temperature.  But [[User:Tohline/SSC/Structure/IsothermalSphere#Isothermal_Sphere|mathematical models of isothermal spheres [*] ]] are relatively cumbersome to analyze because they extend to infinity, they are dynamically unstable, and they are not describable in terms of analytic functions.  In such astrophysical contexts, we have sometimes found it advantageous to employ an $n=5$ polytrope instead of an isothermal sphere.  As in the isothermal case, an [[User:Tohline/SSC/Structure/Polytropes#n_.3D_5_Polytrope|isolated $n=5$ polytrope [*] ]] is infinite in extent and is dynamically unstable, but it is less cumbersome to analyze because its structure can be described by closed-form analytic expressions. + ==Isothermal Spheres== + Isothermal spheres (polytropes with index $n=\infty$) are discussed in a wide variety of astrophysical contexts because it is not uncommon for physical conditions to conspire to create an extended volume throughout which a configuration exhibits uniform temperature.  But [[User:Tohline/SSC/Structure/IsothermalSphere#Isothermal_Sphere|mathematical models of isothermal spheres [*] ]] are relatively cumbersome to analyze because they extend to infinity, they are dynamically unstable, and they are not describable in terms of analytic functions.  In such astrophysical contexts, we have sometimes found it advantageous to employ an $n=5$ polytrope instead of an isothermal sphere.  An [[User:Tohline/SSC/Structure/Polytropes#n_.3D_5_Polytrope|isolated $n=5$ polytrope [*] ]] can serve as an effective surrogate for an isothermal sphere because it is infinite in extent and is dynamically unstable, but it is less cumbersome to analyze because its structure can be described by closed-form analytic expressions. - '''BOUNDED ISOTHERMAL SPHERE & BONNOR-EBERT MASS:'''  In the mid-1950s, [http://adsabs.harvard.edu/abs/1955ZA.....37..217E Ebert] (1955) and [http://adsabs.harvard.edu/abs/1956MNRAS.116..351B Bonnor] (1956) independently realized that an isothermal gas cloud can be stabilized by embedding it in a hot, tenuous external medium.  The relevant mathematical model is constructed by simply chopping off the isothermal sphere at some finite radius and imposing an externally applied pressure, $P_e$, that is equal to the pressure of the isothermal gas at the specified edge of the truncated sphere.  But for a given mass and temperature, there is a value of $P_e$ below which the truncated isothermal sphere is dynamically unstable, like its isolated and untruncated counterpart.  Viewed another way, for a given temperature and $P_e$, a bounded isothermal sphere will be dynamically stable only if its mass is below a critical value, $M_\mathrm{BE}$.  This limiting mass is often referred to as the Bonnor-Ebert mass.  It appears most frequently in the astrophysics literature in discussions of star formation because that is the arena in which both Bonnor and Ebert were conducting research when they made their discoveries. + ==Bounded Isothermal Sphere & Bonnor-Ebert Mass== + In the mid-1950s, [http://adsabs.harvard.edu/abs/1955ZA.....37..217E Ebert] (1955) and [http://adsabs.harvard.edu/abs/1956MNRAS.116..351B Bonnor] (1956) independently realized that an isothermal gas cloud can be stabilized by embedding it in a hot, tenuous external medium.  The relevant mathematical model is constructed by simply chopping off the isothermal sphere at some finite radius and imposing an externally applied pressure, $P_e$, that is equal to the pressure of the isothermal gas at the specified edge of the truncated sphere.  But for a given mass and temperature, there is a value of $P_e$ below which the truncated isothermal sphere is dynamically unstable, like its isolated and untruncated counterpart.  Viewed another way, for a given temperature and $P_e$, a bounded isothermal sphere will be dynamically stable only if its mass is below a critical value, $M_\mathrm{BE}$.  This limiting mass is often referred to as the Bonnor-Ebert mass.  It appears most frequently in the astrophysics literature in discussions of star formation because that is the arena in which both Bonnor and Ebert were conducting research when they made their discoveries. - '''SCHÖNBERG-CHANDRASEKHAR MASS:'''  In the early 1940s, [http://adsabs.harvard.edu/abs/1942ApJ....96..161S Schönberg & Chandrasekhar (1942)] discovered that a star with an isothermal core will become unstable if the fractional mass of the core is above some limiting value.  They discovered this by constructing models that are now commonly referred to as bipolytropes, that is, models in which the star's core is describe + ==Schöberg-Chandrasekhar Mass== + In the early 1940s, [http://adsabs.harvard.edu/abs/1942ApJ....96..161S Schönberg & Chandrasekhar (1942)] discovered that a star with an isothermal core will become unstable if the fractional mass of the core is above some limiting value.  They discovered this by constructing models that are now commonly referred to as bipolytropes, that is, models in which the star's core is describe + + + =Work in Progress= ==Virial Equilibrium== ==Virial Equilibrium== By examining the [[User:Tohline/SphericallySymmetricConfigurations/Virial#P-V_Diagram|virial equilibrium of nonrotating, spherically symmetric configurations that are embedded in an external medium]] of pressure $P_e$, one can begin to appreciate that there is a mass above which no equilibrium exists if the effective adiabatic exponent of the gas is $\gamma_g < 4/3$.  Assuming uniform density and uniform specific entropy configurations for simplicity, there is an analytic expression for this limiting mass and for the equilibrium radius of that limiting configuration.  Specifically, By examining the [[User:Tohline/SphericallySymmetricConfigurations/Virial#P-V_Diagram|virial equilibrium of nonrotating, spherically symmetric configurations that are embedded in an external medium]] of pressure $P_e$, one can begin to appreciate that there is a mass above which no equilibrium exists if the effective adiabatic exponent of the gas is $\gamma_g < 4/3$.  Assuming uniform density and uniform specific entropy configurations for simplicity, there is an analytic expression for this limiting mass and for the equilibrium radius of that limiting configuration.  Specifically,

# Mass Upper Limits

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## Preamble

Spherically symmetric, self-gravitating, equilibrium configurations can be constructed from gases exhibiting a wide variety of degrees of compressibility. When examining how the internal structure of such configurations varies with compressibility, or when examining the relative stability of such structures, it can be instructive to construct models using a polytropic equation of state because the degree of compressibility can be adjusted by simply changing the value of the polytropic index, $~n$, across the range, $0 \le n \le \infty$. (Alternatively, one can vary the effective adiabatic exponent of the gas, γg = 1 + 1 / n.) In particular, n = 0 ($\gamma_g = \infty$) represents a hard equation of state and describes an incompressible configuration, while $n = \infty$ (γg = 1) represents an isothermal and extremely soft equation of state.

## Isolated Polytropes

Isolated polytropic spheres exhibit three attributes that are especially key in the context of our present discussion:

1. The equilibrium structure is dynamically stable if n < 3.
2. The equilibrium structure has a finite radius if n < 5.
3. The equilibrium structure can be described in terms of closed-form analytic expressions for n = 0, n = 1, and n = 5.

## Isothermal Spheres

Isothermal spheres (polytropes with index $n=\infty$) are discussed in a wide variety of astrophysical contexts because it is not uncommon for physical conditions to conspire to create an extended volume throughout which a configuration exhibits uniform temperature. But mathematical models of isothermal spheres [*] are relatively cumbersome to analyze because they extend to infinity, they are dynamically unstable, and they are not describable in terms of analytic functions. In such astrophysical contexts, we have sometimes found it advantageous to employ an n = 5 polytrope instead of an isothermal sphere. An isolated n = 5 polytrope [*] can serve as an effective surrogate for an isothermal sphere because it is infinite in extent and is dynamically unstable, but it is less cumbersome to analyze because its structure can be described by closed-form analytic expressions.

## Bounded Isothermal Sphere & Bonnor-Ebert Mass

In the mid-1950s, Ebert (1955) and Bonnor (1956) independently realized that an isothermal gas cloud can be stabilized by embedding it in a hot, tenuous external medium. The relevant mathematical model is constructed by simply chopping off the isothermal sphere at some finite radius and imposing an externally applied pressure, Pe, that is equal to the pressure of the isothermal gas at the specified edge of the truncated sphere. But for a given mass and temperature, there is a value of Pe below which the truncated isothermal sphere is dynamically unstable, like its isolated and untruncated counterpart. Viewed another way, for a given temperature and Pe, a bounded isothermal sphere will be dynamically stable only if its mass is below a critical value, MBE. This limiting mass is often referred to as the Bonnor-Ebert mass. It appears most frequently in the astrophysics literature in discussions of star formation because that is the arena in which both Bonnor and Ebert were conducting research when they made their discoveries.

## Schöberg-Chandrasekhar Mass

In the early 1940s, Schönberg & Chandrasekhar (1942) discovered that a star with an isothermal core will become unstable if the fractional mass of the core is above some limiting value. They discovered this by constructing models that are now commonly referred to as bipolytropes, that is, models in which the star's core is describe

# Work in Progress

## Virial Equilibrium

By examining the virial equilibrium of nonrotating, spherically symmetric configurations that are embedded in an external medium of pressure Pe, one can begin to appreciate that there is a mass above which no equilibrium exists if the effective adiabatic exponent of the gas is γg < 4 / 3. Assuming uniform density and uniform specific entropy configurations for simplicity, there is an analytic expression for this limiting mass and for the equilibrium radius of that limiting configuration. Specifically,

$M_\mathrm{max}^2 = \frac{3\cdot 5}{2^2 \pi} \biggl(\frac{4}{3\gamma_g} - 1 \biggr) \biggl[ \frac{3 \gamma_g}{4} \biggr] ^{4/(4-3\gamma_g)} \biggl[ \frac{\bar{c_s}^8}{G^3 P_e} \biggr] \, ,$

and,

$R_\mathrm{eq}^2 = \frac{1}{5^2} \biggl[ \frac{4}{3\gamma_g} \biggr]^{2/(4-3\gamma_g)} \biggl[ \frac{GM_\mathrm{max}}{\bar{c_s}^2 } \biggr]^2 = \frac{3}{2^2 \cdot 5\pi} \biggl[ \frac{3\gamma_g}{4} \biggr]^{2/(4-3\gamma_g)} \biggl( \frac{4}{3\gamma_g} - 1 \biggr) \biggl[ \frac{\bar{c_s}^4 }{GP_e} \biggr] \, .$

## Summary of Analytic Results

 Example:   Specify cs, Pe and n (or η = 1 + 1 / n) $~n$ η Mmax Req αvirial αDHB scale αvirial αDHB scale $\infty$ 1 $\biggr(\frac{3^4 \cdot 5^3}{2^{10} \pi}\biggr)^{1/2}$ $\approx 1.77408$ -- $\biggl[ \frac{c_s^8}{G^3 P_e}\biggr]^{1/2}$ $\biggl(\frac{3^2 \cdot 5}{2^6\pi} \bigg)^{1/2}$ $\approx 0.47309$ -- $\biggl( \frac{c_s^4}{GP_e} \biggr)^{1/2}$ 5 6/5 $\biggl( \frac{3^{19}}{2^{12}\cdot 5^9 \pi} \biggr)^{1/2}$ $\approx 0.21505$ $\biggl( \frac{1}{2} \biggr)^{3/10} \biggl( \frac{3^7}{2^8 \pi} \biggr)^{1/2}$ $\approx 1.33943$ $\biggl[ \frac{\bar{c_s}^8}{G^3 P_e}\biggr]^{1/2}$ $\frac{3^{9}}{2^{7}\cdot 5^{6} \pi}$ $\approx 0.0031326$ -- $\biggl( \frac{\bar{c_s}^4}{GP_e} \biggr)^{1/2}$

# Bonnor-Ebert Limiting Mass

Here we will refer to the original works of Ebert (1955, ZA, 37, 217) and Bonnor (1956, MNRAS, 116, 351) .

In his study of the "global gravitational stability [of] one-dimensional polytropes," Whitworth (1981, MNRAS, 195, 967) normalizes (or "references") various derived mathematical expressions for configuration radii, R, and for the pressure exerted by an external bounding medium, Pex, to quantities he refers to as, respectively, Rrf and Prf. The paragraph from his paper in which these two reference quantities are defined is shown here:

In order to map Whitworth's terminology to ours, we note, first, that he uses M0 to represent the spherical configuration's total mass, which we refer to simply as M; and his parameter η is related to our $~n$ via the relation,

$\eta = 1 + \frac{1}{n} \, .$

Hence, Whitworth writes the polytropic equation of state as,

$P = K_\eta \rho^\eta \, ,$

whereas, using our standard notation, this same key relation is written as,

$~P = K_\mathrm{n} \rho^{1+1/n}$ ;

and his parameter Kη is identical to our $~K_\mathrm{n}$.

According to the second (bottom) expression identified by the red outlined box drawn above,

$P_\mathrm{rf} = \frac{3^4 5^3}{2^{10} \pi} \biggl( \frac{K_1^4}{G^3 M^2} \biggr) \, ,$

and inverting the expression inside the green outlined box gives,

$K_1 = \biggl[ K_n (4 P_\mathrm{rf})^{\eta - 1} \biggr]^{1/\eta} \, .$

Hence,

$P_\mathrm{rf} = \frac{3^4 5^3}{2^{10} \pi} \biggl( \frac{1}{G^3 M^2} \biggr)\biggl[ K_n (4 P_\mathrm{rf})^{\eta - 1} \biggr]^{4/\eta} \, ,$

or, gathering all factors of Prf to the left-hand side,

$P_\mathrm{rf}^{(4-3\eta)} = 2^{-2(4+\eta)} \biggl( \frac{3^4 5^3}{\pi} \biggr)^\eta \biggl[ \frac{K_n^4}{G^{3\eta} M^{2\eta}} \biggr] \, .$

Analogously, according to the first (top) expression identified inside the red outlined box,

$R_\mathrm{rf} = \frac{2^2 GM}{3\cdot 5 K_1} = 2^{2/\eta} \biggl( \frac{GM}{3\cdot 5}\biggr) K_n^{-1/\eta} P_\mathrm{rf}^{(1-\eta)/\eta} ~~~~\Rightarrow~~~~ R_\mathrm{rf}^\eta = \frac{2^{2}}{K_n} \biggl( \frac{GM}{3\cdot 5}\biggr)^\eta P_\mathrm{rf}^{(1-\eta)} \, .$

# Related Wikipedia Discussions

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