With the convention used here, a Stokes line is a ray where competing exponential actions have the same imaginary part. A subdominant exponential can acquire a rapidly changing multiplier there. For $e^{\pm x^2/\epsilon }$, these rays have $\arg x=m\pi /2$. The anti-Stokes lines instead have equal exponential magnitudes.

With the equal-magnitude convention, an anti-Stokes line has $\mathrm{Re}(S_1-S_2)=0$ for competing exponentials. For $e^{\pm x^2/\epsilon }$, its rays are $\arg x=\pi /4+m\pi /2$. Crossing one interchanges which exponential grows. Some literature interchanges the names, so the action-based convention should always be specified.

For $z=2n+i\mu \sqrt{n}+O(1)$, the factor $t^{2n}{e}^{z(1-t)}$ has saddle scale $t=1+u/\sqrt{n}$ and becomes $e^{-u^2-i\mu u}$. If the amplitude has a simple pole there, its leading local integral contains $e^{-u^2-i\mu u}/u$. Differentiating its Cauchy principal value with respect to $\mu$ reduces it to a Gaussian integral, producing an error function. The nonsingular amplitude $(1+t{)}^{-1}$ gives $\frac{1}{2}\sqrt{\pi /n}{e}^{-{\mu }^2/4}$.

The symmetric first moment of the traction exerted by a rigid inclusion, with the surface normal pointing into the inclusion. Its contraction with the background rate-of-strain tensor is the strain-related contribution to extra dissipation due to a rigid inclusion. A trace-free version is equivalent for an incompressible flow. This tensor is distinct from the axisymmetric force-dipole flow used as a far-field singularity.

Bodies immersed in one fluid affect each other through the flow that each produces. At large separation, a force-driven sphere generates a Stokeslet that modifies the other sphere's velocity. These interactions produce off-diagonal entries in the hydrodynamic mobility matrix.

A slender low-viscosity fluid column inside a much more viscous fluid rises under a density contrast. At negligible inertia, its axial flux is approximately $q=\pi a^4(\Delta \rho g-P_{i,z})/(8{\mu }_i)$. Radial expansion strains the outer fluid and gives the normal-stress pressure difference $P_i-P_o=2{\mu }_o{a}_t/a$ at leading order when inner normal viscous stress and interfacial tension are negligible. Coupling that pressure to volume conservation yields the conduit equation. This two-fluid viscous model differs from a turbulent entraining thermal plume.

A stock model whose price can be killed or reduced at a default time. The zero-recovery default model combines continuous price dynamics before default with a jump to zero.

A discrete-time stochastic process has bounded increments if a single deterministic $C<\infty$ satisfies $|X_{n+1}-X_n|\le C$ almost surely for every $n$. The constant is uniform in time and sample outcome; bounding each increment by a different constant does not suffice. This property applies to differences, so it does not require the process itself to stay bounded.

For a stochastic process $X_t$ and a fixed sample space outcome $\omega$, its sample path is the function $t\mapsto X_t(\omega )$. One row of observed fair bits, one trajectory of a random walk, or one staircase of Poisson process arrival counts is a sample path. Fixing $t$ instead gives a random variable ranging across possible histories. The same two-variable object $X_t(\omega )$ therefore describes both observations across outcomes and evolution within an outcome.

Statements about continuous, càdlàg or Hölder continuous function sample paths concern these functions simultaneously in time. They differ from assertions made at each fixed time separately. For example, take $U$ with a uniform distribution on $[0,1]$ and set $X_t=1_{\{t=U\}}$. For every fixed $t$, $X_t=0$ almost surely, but every sample path has a single spike. Thus this process is a modification of a stochastic process that is identically zero, while the two are not indistinguishable stochastic processes. Continuity of both sets of sample paths rules out this discrepancy: equality at a countable dense set then implies equality everywhere. This is why path regularity is an additional part of constructing a continuous-time model.

The raw same-observable time correlation $C(t,\tau )=⟨X(t)X(t+\tau )⟩$. For a Stationary process with finite second moments, it equals $⟨X{⟩}^2+\mathrm{Cov}(X(t),X(t+\tau ))$. Its connected part is the autocovariance; unlike the normalized statistical autocorrelation, no division by the variance is imposed. The distinction matters for fluctuations around a nonzero equilibrium point.

A temperature versus density diagram separates classical, degenerate, nonrelativistic and relativistic stellar-matter limits. The electron degeneracy boundary is $T\sim T_F(\rho )$, with slope $2/3$ at low density and $1/3$ at high density. Thermal relativity is controlled by $k_{B}T\sim m_e{c}^2$, while relativity of degenerate electrons is controlled by $p_F\sim m_{e}c$. Photon pressure and thermal electron-positron pairs add further crossovers; all depend on composition and are gradual, not thermodynamic phase boundaries.

In stellar structure, stellar thermal equilibrium means that internal heating balances the surface luminosity and any energy loss through neutrinos, with negligible secular storage of thermal energy on the timescale considered. A star can satisfy this balance while its interior is hotter than its surface and heat flows outward. It is therefore different from equal-temperature thermal equilibrium between systems. Neglecting neutrino losses, the balance is $L_{\mathrm{surface}}=L_{\mathrm{heat}}$. In a slowly mass-losing donor star, a stellar thermal equilibrium radius sequence applies only when thermal adjustment is fast compared with the mass-loss timescale; rapid loss instead requires the adiabatic stellar radius response exponent.

The gas-pressure fraction describes the relative support by an ideal gas and radiation pressure. For monatomic gas plus equilibrium radiation, it determines the adiabatic exponents of a monatomic gas-radiation mixture. Its value at the unperturbed state may be specified, but it generally changes during an isentropic parcel displacement.

For a slowly evolving spherical star, the Lagrangian energy equation separates nuclear rest-mass release, stellar neutrino energy loss, thermal storage and pressure-volume work. Specific internal energy and nuclear energy conventions must avoid counting nuclear mass changes or escaping reaction-neutrino energy twice. At fixed composition the gravothermal stellar energy generation is $-TDS/Dt$. Energy transport processes redistribute heat rather than creating it.

An imposed spherical profile $\rho ={\rho }_c(1-r/R)$ gives $M=\pi {\rho }_c{R}^3/3$ in hydrostatic equilibrium. For an ideal gas with zero surface pressure, $P=P_c(1-x{)}^2(1+2x-9x^2/5)$ and $T=T_c(1-x)(1+2x-9x^2/5)$, where $x=r/R$, $P_c=5G{M}^2/(4\pi R^4)$ and $T_c=5GM/(12\mathcal{R}R)$. The density has a central cusp and temperature initially rises outward, so positive outward radiative diffusion with regular positive nuclear heating cannot satisfy this imposed profile near the centre.

A uniform-density spherical hydrostatic equilibrium model with zero surface pressure has $P(r)=2\pi G{\rho }^2(R^2-r^2)/3$, $M=4\pi \rho R^3/3$ and gravitational energy $-3G{M}^2/(5R)$. A perfect gas can have this mechanical structure with varying temperature and entropy; the imposed equilibrium density does not require incompressible perturbations.

If a population forms at a constant number rate for duration $T$ and every birth object is retained in the census, its present ages are uniform on $[0,T]$. A fixed initial mass function with finite mean birth mass also makes constant mass-based star formation rate equivalent to constant number rate.

Given a birth-mass distribution and age distribution, an evolutionary state occupies the mass-age set where its starting lifetime has elapsed and its ending lifetime has not. The state number fraction is the probability measure of this set. Initial mass function tail coordinates and the probability integral transform can turn these weighted regions into simple geometric areas.

The fluid displacement $\mathbf{\xi }$ determines Eulerian density through $\delta \rho =-\nabla \cdot (\rho \mathbf{\xi })$. The pressure law is $\delta p+\mathbf{\xi }\cdot \nabla p=-\gamma p\nabla \cdot \mathbf{\xi }$. The linear momentum equation must include both $-\delta \rho \nabla \Phi$ and $-\rho \nabla \delta \Phi$, with ${\nabla }^2\delta \Phi =4\pi G\delta \rho$. Retaining the last term incorporates perturbed self-gravity.

For a static spherical star, a fluid displacement $\mathbf{\xi }{e}^{-i\omega t}$ gives ${\rho }_1=-\nabla \cdot (\rho \mathbf{\xi })$, $-{\omega }^2\rho \mathbf{\xi }=-\nabla p_1-{\rho }_1\nabla \Phi -\rho \nabla {\varphi }_1$ and ${\nabla }^2{\varphi }_1=4\pi G{\rho }_1$. The adiabatic relation is $\Delta P=-{\Gamma }_{1}P\nabla \cdot \mathbf{\xi }$. Centre regularity and surface mechanical/gravitational boundary conditions select discrete eigenvalues in a conservative model.