Stellar evolution

Stellar evolution is the change of a star's structure and composition as nuclear reactions, energy transport, mass loss, and interactions proceed.

After substantial envelope loss on the Asymptotic giant branch, a low- or intermediate-mass star's exposed core heats while its radius shrinks, often at approximately constant luminosity during continued shell burning. It may ionize the expelled envelope as a planetary nebula. After the remaining nuclear fuel is depleted, its luminosity decreases and it joins a white dwarf cooling track. Evolutionary tracks on the Hertzsprung-Russell diagram connect the central-star temperature and luminosity with core mass and elapsed time.

The time for stellar nuclear fusion to consume the accessible fuel at characteristic luminosity $L$. If a fraction $f$ of a star's mass $M$ supplies fuel abundance $X$ and released energy per unit fuel mass $E_H$, hydrogen burning gives $t_{\mathrm{nuc}}\sim fXM{E}_H/L$. Accessible core fuel, rather than the entire star, sets the lifetime unless mixing is complete. This is longer than a structural Kelvin-Helmholtz cooling time when the nuclear reservoir greatly exceeds the gravitational contraction reservoir.

A stellar remnant is the object left after a star finishes its principal nuclear-burning evolution, such as a white dwarf, neutron star or black hole. Stellar remnants retain mass while often contributing little optical light, affecting an old stellar population's mass-to-light ratio.

The stellar age is the elapsed time since a star formed. A constant number formation rate on a finite interval gives a uniform distribution of present ages if every formed object remains counted, including stellar remnants. An age distribution inferred only from surviving main sequence stars generally differs because their lifetimes depend on mass.

The Zanstra method estimates a central star's effective temperature by comparing its observed nonionizing continuum with the absorbed ionizing-photon rate inferred from nebular recombination lines. If ${\Phi }_{E_0}(T)={\int }_{E_0/h}^{\infty }{B}_{\nu }(T)/(h\nu )d\nu$, then $[F_{\ell }/(h{\nu }_{\ell }{F}_{{\nu }_o,\ast })][\beta /{\alpha }_{\ell }^{\mathrm{eff}}]={\Phi }_{E_0}(T_s)/B_{{\nu }_o}(T_s)$ when the line traces the full appropriate ionizing budget. Both the distance and stellar radius cancel. Density inhomogeneity also cancels from the integrated recombination estimate when the coefficient ratio is known; photon leakage, dust and mixed-species transfer can invalidate the simple budget.

A carbon-oxygen core is the product of core helium burning. After central helium exhaustion it has little hydrogen or helium, while helium-rich and hydrogen-rich layers surround it.

A blue loop is an excursion to higher effective temperature, followed by a return redward, during core helium burning in some evolved stars. Its extent depends on core and envelope structure, composition gradients and mixing. Only loops reaching the Cepheid instability strip produce the associated additional Cepheid variable crossings.

The Hertzsprung gap is a sparsely populated region of the Hertzsprung-Russell diagram traversed rapidly by many stars evolving away from the main sequence toward the red giant region.

Kelvin-Helmholtz contraction releases gravitational binding energy as a self-gravitating body cools and shrinks. It powers young stars and giant planets before nuclear burning or degeneracy dominates their thermal evolution.

For a monatomic gas obeying the ideal gas law, the specific internal energy is $e=3P/(2\rho )$. With no stellar nuclear fusion heating, the first law of thermodynamics makes the effective stellar energy-generation rate $ϵ=-{\partial }_{t}e-P{\partial }_t(1/\rho )$, with the derivatives taken at fixed enclosed mass. Differentiating gives $ϵ=-3{\partial }_{t}P/(2\rho )+5P{\partial }_t\rho /(2{\rho }^2)$. In a fixed-mass stellar homology Kelvin-Helmholtz contraction, $\rho \propto R^{-3}$ and $P\propto R^{-4}$ at each mass label. Hence $ϵ=-3P\dot{R}/(2\rho R)>0$ when $\dot{R}<0$. This includes internal energy storage as well as pressure work; using the gravitational potential-energy decrease alone would overestimate the outgoing luminosity.

The intrinsic planetary luminosity $L_{\mathrm{int}}$ is the power supplied by interior cooling, contraction, nuclear reactions if present, or internal heating, excluding reflected or passively reradiated stellar light. For fixed mass and no accretion or nuclear source, global conservation of energy gives $L_{\mathrm{int}}=Q_{\mathrm{deep}}-d(U+\Omega )/dt$, where $Q_{\mathrm{deep}}$ is any additional power deposited in the interior, such as tidal heating or Joule heating. For $Q_{\mathrm{deep}}=0$ and a gas-supported quasi-static sphere of fixed specific-heat ratio $\gamma$ and negligible surface pressure, the stellar virial theorem gives $\Omega =-3(\gamma -1)U$, henceFor a dynamically stable gas-supported model take $\gamma >4/3$. For $\gamma =5/3$ this is half the gravitational release. Degenerate giants require a different internal energy reservoir; the half-factor is not universal.

With negligible nuclear or external heating, the stellar structure equations give $L_{\mathrm{int}}=-{\int }_0^{M}T\partial s/\partial tdm$. A nearly isentropic convective interior therefore has $\dot{s}=-L_{\mathrm{int}}/\int Tdm<0$ whenever it radiates. Decreasing specific entropy does not require temperature at every depth to decrease immediately: contraction of an ideal gas can raise its central temperature while total energy is lost.

If otherwise comparable isolated planets of one mass follow $L_{\mathrm{int}}=C{t}^{-\alpha }$, their intrinsic luminosity ratio is $L_{\mathrm{int}}(t_1)/L_{\mathrm{int}}(t_2)=(t_2/t_1{)}^{\alpha }$. The normalization depends on composition and initial entropy. Strong stellar irradiation, deep heating, or accretion can invalidate extrapolating the same cooling law.

The Kelvin-Helmholtz cooling time estimates how long a self-gravitating body can radiate a reservoir of energy of order its gravitational binding energy. It uses intrinsic luminosity $L_{\mathrm{int}}$, excluding reradiated stellar light. At fixed radius and structure, this reservoir scales as $M^2$; the approximation changes when degeneracy or other internal physics changes the accessible energy.

A pre-main-sequence star contracts toward the main sequence before sustained core hydrogen burning supplies its luminosity. Convection, accretion, and deuterium burning shape its track across the Hertzsprung-Russell diagram.

A pre-main-sequence star with a developing radiative core can contract along a relatively horizontal path toward higher effective temperature in the Hertzsprung-Russell diagram. This differs from the nearly vertical fully convective Hayashi track.

A Hayashi track is the nearly vertical path of a fully convective pre-main-sequence star on the Hertzsprung-Russell diagram. Photospheric opacity fixes a weak relation between luminosity and effective temperature.

Match a fully convective stellar polytrope of index $3/2$ to the grey-atmosphere convection onset with eighth-power temperature opacity. At the fixed matching optical depth, $T_b\propto T_e$ and $P_b\propto g^{2/3}{T}_e^{-16/3}$. The ideal-gas adiabat $P=K{\rho }^{5/3}$ gives $K\propto T_b^{5/3}{P}_b^{-2/3}\propto T_e^{47/9}{g}^{-4/9}$. The polytropic mass-radius relation gives $K\propto R{M}^{1/3}$; with $g\propto M/R^2$, this yields $T_e^{47}\propto R{M}^7$. The Stefan–Boltzmann law then gives the displayed luminosity. Composition and opacity normalization are held fixed between models.

Match an efficient, monatomic gas, fully convective star to the grey-atmosphere convection onset with thirteenth-power opacity. Its convective relation is $P=K_T{T}^{5/2}$. At the matching temperature $T_b\propto T_e$, the stellar atmosphere gives $K_T^2\propto M{L}^{-1}{T}_e^{-13}\propto M{R}^{13/2}{L}^{-17/4}$ using the Stefan–Boltzmann law. The ideal-gas relation converts this to $P=K_{\rho }{\rho }^{5/3}$ with $K_{\rho }=(\mathcal{R}/\mu {)}^{5/3}{K}_T^{-2/3}$. The polytropic mass-radius relation at index $3/2$ gives $K_{\rho }\propto G{M}^{1/3}R$, hence $K_T^2\propto M^{-1}{R}^{-3}$. Equating the two expressions proves $L^{17/4}\propto M^2{R}^{19/2}$, or $L\propto M^{8/17}{R}^{38/17}$. The entropy coefficient is constant throughout the convective envelope, but can change between models.

For a monatomic fully convective ideal gas star, uniform specific entropy gives $P=K{T}^{5/2}$ with $K\propto {\mu }^{-5/2}{M}^{-1/2}{R}^{-3/2}$. Combining this with the stellar surface boundary condition $\kappa P=2g/3$ and $\kappa ={\kappa }_0{\rho }^a{T}^b$ givesAll composition comparisons here hold ${\kappa }_0$ fixed. For $a=1,b=4$, $T_e^8\propto {\mu }^4{M}^{2}R$, and the Stefan–Boltzmann law gives $L\propto {\mu }^{2}M{R}^{5/2}$. The stellar virial theorem and Kelvin-Helmholtz contraction then give $R^{-7/2}-R_0^{-7/2}\propto {\mu }^2{M}^{-1}(t-t_0)$, with a positive coefficient. Once the initial-radius term is negligible, $R\propto {\mu }^{-4/7}{M}^{2/7}(t-t_0{)}^{-2/7}$.

For a fixed-mass, fully convective monatomic gas protostar, write $P=K{T}^{5/2}$ and $K=K_0{\mu }^{-5/2}{M}^{-1/2}{R}^{-3/2}$. Matching the photosphere to $\kappa P=2g/3$ with $\kappa ={\kappa }_0\rho T^4$ gives $T_e^8=2G\mathcal{R}{\mu }^4{M}^{2}R/(3{\kappa }_0{K}_0^2)$. The Stefan–Boltzmann law gives $L=\pi ac[2G\mathcal{R}/(3{\kappa }_0){]}^{1/2}{K}_0^{-1}{\mu }^{2}M{R}^{5/2}$. The stellar virial theorem gives total energy $-3G{M}^2/(7R)$, so $L=-3G{M}^2\dot{R}/(7R^2)$. Integrating proves $C=(49\pi ac/6)[2\mathcal{R}/(3{\kappa }_{0}G){]}^{1/2}{K}_0^{-1}$. This holds before stellar nuclear fusion, changing opacity or electron degeneracy pressure invalidates the assumed contraction law.

Match an $n=3/2$ stellar polytrope to the grey-atmosphere convection onset with seventeenth-power opacity. At each instant the convective entropy coefficient obeys $P=K{T}^{5/2}$, $K\propto M^{-1/2}{R}^{-3/2}$ at fixed composition, while the boundary gives $K\propto g^{1/2}{T}_e^{-21/2}$. Thus $T_e^{21/2}\propto M{R}^{1/2}$ and the first displayed luminosity law follows. The stellar virial theorem for a monatomic gas gives total energy $-3G{M}^2/(7R)$, so Kelvin-Helmholtz contraction has $\dot{R}=-7L{R}^2/(3G{M}^2)$. Integrating at fixed mass gives $R^{-67/21}-R_0^{-67/21}\propto t$. The second displayed relation holds after the initial-radius term is negligible. The entropy coefficient is spatially constant but changes during contraction.

For convective contraction with power-law opacity $\kappa \propto \rho T^4$ at fixed stellar mass and opacity normalization, $T_e\propto {\mu }^{1/2}{R}^{1/8}$ and $L\propto {\mu }^2{R}^{5/2}$. Eliminating $R$ gives $L\propto {\mu }^{-8}{T}_e^{20}$. Thus a larger mean molecular weight moves the track to larger effective temperature at a given luminosity, with $T_e{|}_L\propto {\mu }^{2/5}$. Contraction decreases both temperature and luminosity for this particular opacity exponent.

Matter accreting at rate $\dot{M}$ onto a star of mass $M$ and radius $R$ releases gravitational power of orderup to an efficiency and an inner-boundary correction.

At the temperatures relevant to pre-main-sequence stars, a useful steep power-law approximation to the specific deuterium-burning rate is ${ϵ}_D\propto \rho T^{15}$. Its sensitivity to central temperature makes deuterium burning act as a temporary thermostat.

A subgiant has begun evolving away from the main sequence after central hydrogen exhaustion, increasing its radius before the extended red giant stage.

A red giant is an evolved star with a large radius and low effective temperature. On the first giant branch, a low- or intermediate-mass star is powered primarily by shell hydrogen burning around a helium-rich core.

Use a Kramers radiative-zero envelope around a stellar core and specific heating $ϵ={ϵ}_0\rho T^{67/4}$. Since $\rho \propto C{T}^{13/4}$ and $T\propto M_c/r$, $ϵ\propto C{r}^{-20}$ and $ϵ(1.05R_c)/ϵ(R_c)=1.05^{-20}\simeq 0.377$. The volume heating is $\rho ϵ\propto C^2{T}^{93/4}$. Integrating $4\pi r^2\rho ϵ$ from $R_c$ outwards gives $L\propto C^2{M}_c^{93/4}{R}_c^{-81/4}$, while radiative transport gives $C^2\propto M_c/L$. Therefore $L^2\propto M_c^{97/4}{R}_c^{-81/4}$, or $L\propto M_c^{97/8}{R}_c^{-81/8}$. For a core sequence with fixed $R_c$, the mass exponent is $97/8$. The steep integrand confines most heating to a thin shell; using the exterior constant-luminosity profile throughout it estimates the coefficient rather than resolving the shell structure.

In a thin-shell giant model with a nonrelativistic degenerate core, $R_c\propto M_c^{-1/3}$. For a gas-pressure radiative layer with $\kappa \propto \rho T^{-3}$ and specific energy generation $ϵ\propto \rho T^{10}$, the deep profile has $P\propto T^4$, $\rho \propto T^3$ and $T\propto M_c/r$. Integrating the strongly concentrated shell power gives $L\propto C^2{M}_c^{16}{R}_c^{-13}$, while radiative transport gives $C^2\propto M_c/L$. Therefore $L^2\propto M_c^{17}{R}_c^{-13}$, yielding the displayed exponent. This is a model-specific scaling, not a universal core-mass law; it fails when envelope stripping removes the shell fuel or the confining structure.

A negligible-mass radiative envelope surrounding a core of mass $M_c$, carrying constant luminosity and having constant opacity, obeys $dP/dT=16\pi acG{M}_c{T}^3/(3{\kappa }_{0}L)$. Neglecting the outer pressure integration constant gives $\Lambda =4\pi acG{M}_c/(3{\kappa }_{0}L)$. The gas pressure fraction $\beta$ is constant because $P_{\mathrm{rad}}=a{T}^4/3$, and $\rho =\mu \beta \Lambda T^3/\mathcal{R}$. Hydrostatic balance gives $T=T_s+\mu \beta G{M}_c(1/r-1/R_s)/(4\mathcal{R})$; dropping this outer-boundary constant gives the displayed deep-envelope approximation. A positive gas-pressure fraction requires $L<4\pi cG{M}_c/{\kappa }_0$.

With specific heating $ϵ={ϵ}_0\rho T^{16}$, the constant-opacity radiative-zero giant envelope has $T\propto \mu \beta M_c/r$ and $\rho \propto \mu \beta M_c{T}^3/L$. Therefore the shell power $4\pi r^2\rho ϵ$ is proportional to $(\mu \beta M_c{)}^{24}{L}^{-2}{r}^{-20}$. Integrating from the core radius gives $L^3\propto (\mu \beta M_c{)}^{24}{R}_c^{-19}$. The steep radial dependence justifies a thin-shell estimate dominated by the base. Using the outer radiative profile throughout the emitting layer is an approximation; a resolved shell can change the numerical coefficient. Also $\beta =1-{\kappa }_{0}L/(4\pi cG{M}_c)$, so retaining radiation pressure makes the relation implicit in luminosity.

Stellar rejuvenation occurs when mixing brings fresh hydrogen into the core of a mass-gaining or merged main sequence star, making its remaining nuclear lifetime resemble that of a younger star. The amount of rejuvenation depends on the accretor's internal mixing and prior composition.

A helium star has lost most of its hydrogen-rich envelope and is made predominantly of helium. If massive enough, its core can ignite helium; a lower-mass stripped core instead becomes a helium white dwarf.

An asymptotic-giant-branch star has an inert carbon-oxygen core surrounded by helium- and hydrogen-burning shells and a deep convective envelope. Alternating shell burning produces thermal pulses and dredge-up, while strong mass loss removes the envelope.

The fate of an Asymptotic giant branch star depends on whether its envelope is removed before its core reaches the next ignition or instability threshold. The approximate envelope-loss time is $t_{\mathrm{env}}=M_{\mathrm{env}}/|\dot{M}|$, while the time to add a required core mass $\Delta M_c$ is $t_{\mathrm{grow}}=\Delta M_c/{\dot{M}}_c$. Shell stellar nuclear fusion grows the core; third dredge-up can reduce the net growth. If $t_{\mathrm{env}}\ll t_{\mathrm{grow}}$, shell fuel disappears first and a white dwarf remnant is favored. The opposite ordering permits further core growth but does not by itself specify whether carbon burning is stable, explosive, or followed by further evolution. Both rates change with time, so the estimates guide rather than replace a stellar model.

When the bottom of an Asymptotic giant branch star's convective envelope becomes hot enough for hydrogen burning, the CNO cycle can convert newly dredged-up carbon into nitrogen. Its onset and strength depend on mass, composition and envelope structure; it is not a universal property of all stars at one initial mass.

Stellar dredge-up occurs when a convective envelope penetrates into layers altered by nuclear burning and transports their material to the surface.

First dredge-up occurs as a star ascends the red-giant branch. Its deepening convective envelope exposes material processed by hydrogen burning, especially enhanced helium and nitrogen and depleted carbon.

Second dredge-up occurs after core helium exhaustion in sufficiently massive intermediate-mass stars. The convective envelope moves inward, mixes helium- and hydrogen-burning products to the surface, and reduces the hydrogen-exhausted core mass.

Third dredge-up follows a thermal pulse on the asymptotic giant branch, when the convective envelope reaches into the intershell and can carry freshly made carbon and slow-neutron-capture elements to the surface.

A thermal pulse is a helium-shell flash caused by thermally unstable burning in a thin shell. It briefly drives intershell convection, expands and cools the hydrogen-burning shell, and can be followed by third dredge-up.

The Härm–Schwarzschild instability is the thin-shell thermal instability of helium burning in an asymptotic-giant-branch star. Heating greatly raises the nuclear rate while the geometrically thin shell initially expands too little to provide stabilizing cooling.

A super-AGB star is massive enough to ignite carbon and form a degenerate oxygen-neon core, but initially too light to proceed directly through all advanced burning stages. Its fate depends on competition between core growth and envelope mass loss.

An oxygen-neon-magnesium core is the degenerate product of carbon burning, composed mainly of oxygen and neon with a smaller magnesium abundance. It can become an oxygen-neon white dwarf or collapse after electron captures if it approaches the Chandrasekhar mass.

An electron-capture supernova occurs when electron captures, especially on neon and magnesium nuclei in a degenerate oxygen-neon core, remove pressure support and trigger collapse before ordinary iron-core formation.

A stellar wind is a continuous outflow of gas from a star. Radiation, gas pressure, waves, or magnetic fields can supply its acceleration, and the resulting mass loss can determine a star's final remnant.

An episode of particularly strong mass loss near the end of Asymptotic giant branch evolution, capable of removing the remaining envelope rapidly compared with its prior evolution. Radiation acting on dust can transfer momentum to the gas; pulsations help lift material into cooler dust-forming layers. Envelope removal terminates sustained hydrogen-burning shell and helium-burning shell activity.

A magnetized stellar wind extracts angular momentum through matter transport and Maxwell stress tensor. The Alfvén radius supplies an effective lever arm larger than the stellar radius, with geometry factors absorbed into an averaged lever arm. Dynamo action couples the torque to rotation by supplying the magnetic field. For a simplified radial wind of fixed speed, $R_A^2$ scales as $B_{\ast }^2{R}_{\ast }^4/({\mu }_0{\dot{M}}_w{v}_w)$, so unsaturated $B_{\ast }\propto \Omega$ gives a torque proportional to $-{\Omega }^3$. Magnetic topology and wind thermodynamics can alter this scaling.

If the moment of inertia I changes slowly and wind-driven magnetic braking of a solar-type star has torque $-K{\Omega }^3$, then ${\Omega }^{-2}(t)={\Omega }^{-2}(t_0)+2K(t-t_0)/I$. At sufficiently late times within this unsaturated approximation, rotation scales as $t^{-1/2}$. This empirical main-sequence trend is not a universal law for saturated rapid rotators, changing stellar structure or every magnetic geometry.

Using an axisymmetric hydrodynamic mass flux function, set $s=s(\psi )$, $\ell =R{u}_{\varphi }=\ell (\psi )$ and $\epsilon =\epsilon (\psi )$. The azimuthal vorticity is ${\omega }_{\varphi }=-R\nabla \cdot (\nabla \psi /(\rho R^2))$. The meridional projection of the Crocco form of the unsteady Euler equation gives the displayed equation. The mass density also satisfies the Bernoulli relation $\epsilon =|\nabla \psi {|}^2/(2{\rho }^2{R}^2)+{\ell }^2/(2R^2)+w+\Phi$, so an equation of state is part of closing the system.

A steady spherical polytropic flow escaping a point-mass potential conserves the outward mass loss rate and the Bernoulli function $C=u^2/2+c_s^2/(\gamma -1)-GM/r$. For positive terminal speed, a finite-radius smooth sonic point requires $C>0$ and $1<\gamma <5/3$. Its sound speed obeys $c_{s,c}^2=2C(\gamma -1)/(5-3\gamma )$, and evaluating the conserved flux there selects the transonic mass loss rate. Other parameter limits need separate treatment.

Put $y=\rho r^{4/(\gamma +1)}$ and $\alpha =4(\gamma -1)/(\gamma +1)$. The polytropic stellar wind equations become $F(r)=C{r}^{\alpha }+GM{r}^{\alpha -1}=G(y)=\gamma K{y}^{\gamma -1}/(\gamma -1)+{\dot{M}}^2/(32{\pi }^2{y}^2)$. At a regular sonic point both derivatives vanish. Matching their minima allows the two positive-$y$ branches to join with finite slope; unequal minima either leave a forbidden radial interval or separate the subsonic and supersonic branches. This reduces a singular differential-equation crossing to a geometric minimum comparison.

In an optically thick steady spherical stellar wind of constant mean molecular weight ${\mu }_{\mathrm{mol}}$, let $v$ be the outward speed and $a_g^2=k_{B}T/({\mu }_{\mathrm{mol}}{m}_u)$ the isothermal sound speed squared. Mass conservation gives ${\rho }^{\prime }/\rho =-2/r-v^{\prime }/v$, where a prime denotes differentiation with respect to radius $r$. Use $P_{\mathrm{gas}}=\rho a_g^2$ and radiative diffusion to write $P_{\mathrm{rad}}^{\prime }/\rho =-g\Gamma$, where $g=G{M}_r/r^2$ and $\Gamma =L_r/L_{\mathrm{Edd}}$ uses the local Eddington luminosity $L_{\mathrm{Edd}}=4\pi cG{M}_r/\kappa$ and opacity $\kappa$. Substituting into the momentum equation yields the displayed wind equation. Smooth sonic point crossing requires both its denominator and numerator to vanish. The radiation-pressure gradient already supplies radiative acceleration; adding that force separately would count it twice.

For an optically thick diffusion-controlled wind, put $\beta =P_{\mathrm{gas}}/(P_{\mathrm{gas}}+P_{\mathrm{rad}})$. Diffusion gives $(\log T{)}^{\prime }=-g\Gamma \beta /[4a_g^2(1-\beta )]$. At $\Gamma =1$, the numerator of the gas and radiation pressure stellar wind equation is $2a_g^2/r+g\beta /[4(1-\beta )]>0$ for finite-temperature gas and $0<\beta <1$. A subsonic branch then decelerates and cannot pass a regular sonic point. This statement is specific to subsonically launched winds in this transport model, not to every supersonic solution or every form of stellar-wind driving.

A neutron star is a compact stellar remnant whose gravity compresses roughly a stellar mass into a radius of order ten kilometres. Degenerate matter supplies most of its pressure support, while rapid rotation and intense magnetic fields can power observable emission.

A pulsar is a rotating magnetized neutron star whose directed emission is observed as periodic pulses. Rotation supplies the energy of an ordinary radio pulsar.

A pulsar magnetosphere is the plasma and electromagnetic field surrounding a pulsar. Close to the star, charges tend to screen electric fields parallel to the magnetic field and approximately corotate; near the light cylinder, relativistic motion and field distortion become essential.

In an idealized corotating pulsar magnetosphere, plasma moves with $\mathbf{u}=\mathbf{\Omega }\times \mathbf{r}$. In Gaussian units, ideal magnetohydrodynamics then gives $\mathbf{E}=-\mathbf{u}\times \mathbf{B}/c$.

The Goldreich-Julian density is the charge density needed to support corotation in an idealized pulsar magnetosphere. Where exterior currents vanish,in Gaussian units.

The light cylinder is the notional cylindrical surface on which rigid corotation with angular speed $\Omega$ would reach the speed of light.

The light cylinder radius is $R_{\mathrm{LC}}=c/\Omega$. A nonrelativistic corotation model must fail as cylindrical radius approaches this scale.

A pulsar polar cap is the region on the stellar surface threaded by magnetic field lines that reach the outer magnetosphere rather than closing close to the star.

For an aligned dipole whose last relevant field line reaches equatorial radius $R_e$, the polar-cap opening angle obeys ${\sin }^2{\theta }_{\ast }=R_{\ast }/R_e$. It is approximately $\sqrt{R_{\ast }/R_e}$ for a small cap.

A white dwarf is a compact stellar remnant supported mainly by electron degeneracy pressure.

A carbon-oxygen white dwarf is a white dwarf whose core is predominantly carbon and oxygen, retained after core helium burning when the star loses its envelope without sustained carbon burning.

A helium white dwarf is a white dwarf whose core is primarily helium. Stripping a low-mass evolved star before helium ignition can produce one on a much shorter timescale than isolated stellar evolution.

A cold nonrelativistic white dwarf supported by electron degeneracy pressure has a polytropic equation of state $P=K_{\mathrm{NR}}{\rho }^{5/3}$, so its polytropic index is $3/2$. The Lane-Emden mass formula and polytropic mass-radius relation give $R\propto K_{\mathrm{NR}}{M}^{-1/3}$ at fixed composition. As density rises, relativistic corrections soften the exponent toward $4/3$ and lead toward the Chandrasekhar mass regime.

The Chandrasekhar mass is the limiting mass of a cold nonrotating white dwarf supported by relativistic electron degeneracy pressure, about $1.4M_{\odot }$ for ordinary composition.

An ultrarelativistic degenerate electron gas has $P=K{\rho }^{4/3}$, so it forms an index-three stellar polytrope. Insert $P_c=c_{3}G{M}^2/R^4$ and ${\rho }_c=3a_{3}M/(4\pi R^3)$ into that equation of state: the radius cancels, leaving the displayed mass. For fully ionized helium, $n_e=\rho /(2m_p)$ and $K=(3/(2\pi ){)}^{1/3}hc/(16m_p^{4/3})$. This is the ideal limiting Chandrasekhar mass, rather than a claim that every ordinary white dwarf has that mass.

Mestel's cooling law matches a degenerate isothermal white-dwarf core to a nondegenerate radiative envelope. With Kramers opacity it gives $L/M\propto T_c^{7/2}$ in terms of core temperature.

In a negligible-mass white dwarf envelope, take enclosed mass and luminosity constant, neglect radiation pressure, and use an ideal gas with the Kramers' opacity law. Dividing hydrostatic pressure support equation by radiative diffusion in a star gives $PdP\propto (M/L)T^{15/2}dT$. Neglecting the surface integration constant produces the displayed pressure and mass density laws. It is a deep-envelope approximation, not a finite-temperature photospheric boundary condition.

Integrating the hydrostatic pressure support equation with the radiative-zero white-dwarf envelope law yields $T=4GM(1/r-1/R)/(17\mathcal{R})$. Matching $T=T_c$ at $r_b$ gives $s=R-r_b$ as displayed. A strongly degenerate bulk core has thermal energy small compared with its Fermi energy, making $q\ll 1$ and hence $s\ll R$. The inequality $T_c\gg T_{\mathrm{eff}}$ alone does not establish geometrical thinness.

For fixed mass, degenerate radius and composition, the matching of a white-dwarf envelope to a degenerate core gives $L\propto T_c^{7/2}$, while the geometrical thickness of a white-dwarf radiative envelope gives $s\propto T_c$ to leading thin-layer order. Cooling therefore thins this idealized envelope. At finite thickness the ratio additionally contains $(1+q_2)/(1+q_1)$.

Continuity of temperature, pressure and mass density matches the ideal gas envelope pressure $\mathcal{R}{\rho }_b{T}_c$ to the nonrelativistic electron degeneracy pressure $K_{\mathrm{NR}}{\rho }_b^{5/3}$. Substituting this base mass density into the radiative-zero white-dwarf envelope law gives the luminosity-temperature scaling behind Mestel's cooling law. The matched point is a transition approximation; the cold degeneracy law does not represent finite-temperature corrections there exactly.

A hydrogen-burning white-dwarf envelope is a thin nondegenerate surface layer in which accreted or residual hydrogen can undergo nuclear fusion above a degenerate core.

In a thin nondegenerate hydrogen-burning white-dwarf envelope, fix gravity $g$ and use $P=g\Sigma$, $\rho =B\Sigma /T$. With opacity $\kappa ={\kappa }_0\rho T^{-2}$ and specific heating $ϵ={ϵ}_0\rho T^{15}$, set $y=T^7$, $x={\Sigma }^2/2$. Then $y^{\prime }=AF$ and $F^{\prime }=-{ϵ}_{0}B{y}^2$, where $A=21{\kappa }_{0}B/(4ac)$ and ${\omega }^2=A{ϵ}_{0}B$. Radiative-zero surface and zero core-flux base conditions give $(y^{\prime }{)}^2=(2{\omega }^2/3)(y_b^3-y^3)$. Thus $T_b\propto F_s^{2/21}$ and $F_s=12I^3/(A{\omega }^2{\Sigma }_b^6)$, with $I={\int }_0^1(1-t^3{)}^{-1/2}dt$. At fixed column the specific heating varies as $T^{14}$ while radiative cooling varies as $T^7$, giving thin-shell instability in this approximation.

Thin-shell instability is a thermal runaway of nuclear burning in a geometrically thin layer whose pressure is constrained by overlying weight, weakening expansion-driven cooling.

A classical nova is a thermonuclear runaway in hydrogen-rich material accreted onto a white dwarf, producing a sudden luminous outburst and ejecting part of the envelope.

An idealized classical nova that slowly ejects mass $\delta m$ isotropically from the accretor, compared with the orbital period, but quickly compared with the mass-transfer timescale, expands a circular binary star's Roche lobe. The ejecta carry accretor specific angular momentum $j_1/J=M_2/(M_{1}M)=q/M$. Taking $\delta J/J=-q\delta m/M$, $\delta M_1=-\delta m$ and $\delta M_2=0$ in $J=M_1{M}_2\sqrt{Ga/M}$ proves $\delta a/a=\delta m/M$. The lobe law $R_L\propto a(M_2/M{)}^{1/3}$ then gives $\delta R_L/R_L=4\delta m/(3M)$. A donor whose radius is unchanged moves inside its lobe.

For constant external loss $\Gamma =-\dot{J}>0$ and a donor sequence $R_2\propto M_2^{\zeta }$, the fractional gap closes at rate $2\Gamma /J$, giving $t_d=2J\delta m/(3M\Gamma )$. The binary mass-transfer contact equation gives accumulation time $t_s=J\delta m(\zeta +5/3-2q)/(2M_2\Gamma )$. ThereforeThis first-order result requires $\delta m/M\ll 1$, a positive contact denominator and a fixed donor radius during the eruption. At $\zeta =1$ it reduces to $2q/[(1+q)(4-3q)]$.

The Schönberg-Chandrasekhar limit is the maximum fractional mass of an isothermal inert core that can support a surrounding hydrogen-burning envelope in hydrostatic and thermal equilibrium.