Perturbations and problems of two bodies

The approximate nature of Kepler’s laws

The constraints placed on the force for Kepler’s laws to be derivable from Newton’s laws were that the force must be directed toward a central fixed point and that the force must decrease as the inverse square of the distance. In actuality, however, the Sun, which serves as the source of the major force, is not fixed but experiences small accelerations because of the planets, in accordance with Newton’s second and third laws. Furthermore, the planets attract one another, so that the total force on a planet is not just that due to the Sun; other planets perturb the elliptical motion that would have occurred for a particular planet if that planet had been the only one orbiting an isolated Sun. Kepler’s laws therefore are only approximate. The motion of the Sun itself means that, even when the attractions by other planets are neglected, Kepler’s third law must be replaced by (M + m_i)τ² ∝ a³, where m_i is one of the planetary masses and M is the Sun’s mass. That Kepler’s laws are such good approximations to the actual planetary motions results from the fact that all the planetary masses are very small compared to that of the Sun. The perturbations of the elliptic motion are therefore small, and the coefficient M + m_i ≈ M for all the planetary masses m_i means that Kepler’s third law is very close to being true.

Newton’s second law for a particular mass is a second-order differential equation that must be solved for whatever forces may act on the body if its position as a function of time is to be deduced. The exact solution of this equation, which resulted in a derived trajectory that was an ellipse, parabola, or hyperbola, depended on the assumption that there were only two point particles interacting by the inverse square force. Hence, this “gravitational two-body problem” has an exact solution that reproduces Kepler’s laws. If one or more additional bodies also interact with the original pair through their mutual gravitational interactions, no exact solution for the differential equations of motion of any of the bodies involved can be obtained. As was noted above, however, the motion of a planet is almost elliptical, since all masses involved are small compared to the Sun. It is then convenient to treat the motion of a particular planet as slightly perturbed elliptical motion and to determine the changes in the parameters of the ellipse that result from the small forces as time progresses. It is the elaborate developments of various perturbation theories and their applications to approximate the exact motions of celestial bodies that has occupied celestial mechanicians since Newton’s time.

Perturbations of elliptical motion

So far the following orbital parameters, or elements, have been used to describe elliptical motion: the orbital semimajor axis a, the orbital eccentricity e, and, to specify position in the orbit relative to the perihelion, either the true anomaly f, the eccentric anomaly u, or the mean anomaly l. Three more orbital elements are necessary to orient the ellipse in space, since that orientation will change because of the perturbations. The most commonly chosen of these additional parameters are illustrated in Figure 2, where the reference plane is chosen arbitrarily to be the plane of the ecliptic, which is the plane of Earth’s orbit defined by the path of the Sun on the sky. (For motion of a near-Earth artificial satellite, the most convenient reference plane would be that of Earth’s Equator.) Angle i is the inclination of the orbital plane to the reference plane. The line of nodes is the intersection of the orbit plane with the reference plane, and the ascending node is that point where the planet travels from below the reference plane (south) to above the reference plane (north). The ascending node is described by its angular position measured from a reference point on the ecliptic plane, such as the vernal equinox; the angle Ω is called the longitude of the ascending node. Angle ω (called the argument of perihelion) is the angular distance from the ascending node to the perihelion measured in the orbit plane.

For the two-body problem, all the orbital parameters a, e, i, Ω, and ω are constants. A sixth constant T, the time of perihelion passage (i.e., any date at which the object in orbit was known to be at perihelion), may be used to replace f, u, or l, and the position of the planet in its fixed elliptic orbit can be determined uniquely at subsequent times. These six constants are determined uniquely by the six initial conditions of three components of the position vector and three components of the velocity vector relative to a coordinate system that is fixed with respect to the reference plane. When small perturbations are taken into account, it is convenient to consider the orbit as an instantaneous ellipse whose parameters are defined by the instantaneous values of the position and velocity vectors, since for small perturbations the orbit is approximately an ellipse. In fact, however, perturbations cause the six formerly constant parameters to vary slowly, and the instantaneous perturbed orbit is called an osculating ellipse; that is, the osculating ellipse is that elliptical orbit that would be assumed by the body if all the perturbing forces were suddenly turned off.

First-order differential equations describing the variation of the six orbital parameters can be constructed for each planet or other celestial body from the second-order differential equations that result by equating the mass times the acceleration of a body to the sum of all the forces acting on the body (Newton’s second law). These equations are sometimes called the Lagrange planetary equations after their derivation by the great Italian-French mathematician Joseph-Louis Lagrange (1736–1813). As long as the forces are conservative and do not depend on the velocities—i.e., there is no loss of mechanical energy through such processes as friction—they can be derived from partial derivatives of a function of the spatial coordinates only, called the potential energy, whose magnitude depends on the relative separations of the masses.

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In the case where all the forces are derivable from such potential energy, the total energy of a system of any number of particles—i.e., the kinetic energy plus the potential energy—is constant. The kinetic energy of a single particle is one-half its mass times the square of its velocity, and the total kinetic energy is the sum of such expressions for all the particles being considered. The conservation of energy principle is thus expressed by an equation relating the velocities of all the masses to their positions at any time. The partial derivatives of the potential energy with respect to spatial coordinates are transformed into particle derivatives of a disturbing function with respect to the orbital elements in the Lagrange equations, where the disturbing function vanishes if all bodies perturbing the elliptic motion are removed. Like Newton’s equations of motion, Lagrange’s differential equations are exact, but they can be solved only numerically on a computer or analytically by successive approximations. In the latter process, the disturbing function is represented by a Fourier series, with convergence of the series (successive decrease in size and importance of the terms) depending on the size of the orbital eccentricities and inclinations. Clever changes of variables and other mathematical tricks are used to increase the time span over which the solutions (also represented by series) are good approximations to the real motion. These series solutions usually diverge, but they still represent the actual motions remarkably well for limited periods of time. One of the major triumphs of celestial mechanics using these perturbation techniques was the discovery of Neptune in 1846 from its perturbations of the motion of Uranus.