Nbody problem
In physics, the nbody problem is an ancient, classical problem^{[1]} of predicting the individual motions of a group of celestial objects interacting with each other gravitationally. Solving this problem — from the time of the Greeks and on — has been motivated by the desire to understand the motions of the Sun, Moon, planets and the visible stars. In the 20th century, understanding the dynamics of globular cluster star systems became an important nbody problem.^{[2]} The nbody problem in general relativity is considerably more difficult to solve.
The classical physical problem can be informally stated as: given the quasisteady orbital properties (instantaneous position, velocity and time)^{[3]} of a group of celestial bodies, predict their interactive forces; and consequently, predict their true orbital motions for all future times.^{[4]}
To this purpose the twobody problem has been completely solved and is discussed below; as is the famous Restricted 3Body Problem.^{[5]}
Contents
 History 1
 General formulation 2

Special cases 3
 Twobody problem 3.1
 Threebody problem 3.2
 Planetary problem 3.3
 Central configurations 3.4
 nbody choreography 3.5

Analytic approaches 4
 Power series solution 4.1
 A generalized Sundman global solution 4.2
 Singularities of the nbody problem 4.3

Simulation 5
 Few bodies 5.1
 Many bodies 5.2
 Strong gravitation 5.3
 Other nbody problems 6
 See also 7
 Notes 8
 References 9
 Further reading 10
 External links 11
History
Knowing three orbital positions of a planet's orbit – positions obtained by Sir Isaac Newton (16431727) from astronomer John Flamsteed^{[6]} – Newton was able to produce an equation by straightforward analytical geometry, to predict a planet's motion; i.e., to give its orbital properties: position, orbital diameter, period and orbital velocity.^{[7]} Having done so he and others soon discovered over the course of a few years, those equations of motion did not predict very well or even correctly some orbits^{[8]}^{,};^{[9]} and Newton realized it was because gravitational interactive forces amongst all the planets was affecting all their orbits.
The above discovery goes right to the heart of the matter as to what exactly the nbody problem is physically: as Newton realized, it is not sufficient to just specify the initial position and velocity, or three orbital positions either, to determine a planet's true orbit: the gravitational interactive forces have to be known too. Thus came the awareness and rise of the nbody “problem” in the early 17th century. These gravitational attractive forces do conform to Newton's Laws of Motion and to his Law of Universal Gravitation, but the many multiple (nbody) interactions have historically made any exact solution intractable. Ironically, this conformity led to the wrong approach.
After Newton's time the nbody problem historically was not stated correctly because it did not include a reference to those gravitational interactive forces. Newton does not say it directly but implies in his Principia the nbody problem is unsolvable because of those gravitational interactive forces.^{[10]} Newton said^{[11]} in his Principia, paragraph 21:
And hence it is that the attractive force is found in both bodies. The Sun attracts Jupiter and the other planets, Jupiter attracts its satellites and similarly the satellites act on one another. And although the actions of each of a pair of planets on the other can be distinguished from each other and can be considered as two actions by which each attracts the other, yet inasmuch as they are between the same, two bodies they are not two but a simple operation between two termini. Two bodies can be drawn to each other by the contraction of rope between them. The cause of the action is twofold, namely the disposition of each of the two bodies; the action is likewise twofold, insofar as it is upon two bodies; but insofar as it is between two bodies it is single and one ...
Newton concluded via his 3rd Law that "according to this Law all bodies must attract each other." This last statement, which implies the existence of gravitational interactive forces, is key.
As shown below, the problem also conforms to Jean Le Rond D'Alembert's nonNewtonian 1st and 2nd Principles and to the nonlinear nbody problem algorithm, the latter allowing for a closed form solution for calculating those interactive forces.
The problem of finding the general solution of the nbody problem was considered very important and challenging. Indeed in the late 19th century King Oscar II of Sweden, advised by Gösta MittagLeffler, established a prize for anyone who could find the solution to the problem. The announcement was quite specific:
Given a system of arbitrarily many mass points that attract each according to Newton's law, under the assumption that no two points ever collide, try to find a representation of the coordinates of each point as a series in a variable that is some known function of time and for all of whose values the series converges uniformly.
In case the problem could not be solved, any other important contribution to classical mechanics would then be considered to be prizeworthy. The prize was awarded to Poincaré, even though he did not solve the original problem. (The first version of his contribution even contained a serious error^{[12]}). The version finally printed contained many important ideas which led to the development of chaos theory. The problem as stated originally was finally solved by Karl Fritiof Sundman for n = 3.
General formulation
The nbody problem considers N point masses m_i, i=1,2,\ldots,N in an inertial reference in three dimensional space \mathbb{R}^3 moving under the influence of mutual gravitational attraction. Each mass m_i has a position vector \mathbf{q}_i. Newton's second law says that mass times acceleration m_i d^2\mathbf{q}_i/dt^2 is equal to the sum of the forces on the mass. Newton's law of gravity says that the gravitational force felt on mass m_i by a single mass m_j is given by^{[13]}
 \mathbf{F}_{ij} = \frac{G m_i m_j (\mathbf{q}_j  \mathbf{q}_i)}{\left\ \mathbf{q}_j  \mathbf{q}_i\right\^3},
where G is the gravitational constant and \left\ \mathbf{q}_j  \mathbf{q}_i\right\ is the magnitude of the distance between \mathbf{q}_i and \mathbf{q}_j.
Summing over all masses yields the nbody equations of motion:
 m_i \frac{d^2\mathbf{q}_i}{dt^2} = \sum_{j=1,i \ne j}^N \frac{G m_i m_j (\mathbf{q}_j  \mathbf{q}_i)}{\left\ \mathbf{q}_j  \mathbf{q}_i\right\^3} = \frac{\partial U}{\partial \mathbf{q}_i}
where U is the selfpotential energy
 U = \sum_{1 \le i < j \le N} \frac{G m_i m_j}{\left\ \mathbf{q}_j  \mathbf{q}_i\right\}.
Defining the momentum to be \mathbf{p}_i = m_i d\mathbf{q}_i/dt, Hamilton's equations of motion for the nbody problem become^{[14]}
 \frac{d\mathbf{q}_i}{dt} = \frac{\partial H}{\partial \mathbf{p}_i} \qquad \frac{d\mathbf{p}_i}{dt} = \frac{\partial H}{\partial \mathbf{q}_i},
where the Hamiltonian function is
 H = T + U
and T is the kinetic energy
 T = \sum_{i=1}^N \frac{\left\ \mathbf{p}_i \right\^2}{2m_i}.
Hamilton's equations show that the nbody problem is a system of 6N firstorder differential equations, with 6N initial conditions as 3N initial position coordinates and 3N initial momentum values.
Symmetries in the nbody problem yield global integrals of motion that simplify the problem.^{[15]} Translational symmetry of the problem results in the center of mass
 \mathbf{C} = \frac{\sum_{i=1}^N m_i \mathbf{q}_i}{\sum_{i=1}^N m_i}
moving with constant velocity, so that \mathbf{C} = \mathbf{L}_0 t + \mathbf{C}_0 , where \mathbf{L}_0 is the linear momentum and C_{0} is the initial position. The constants of motion L_{0} and \mathbf{C}_0 represent six integrals of the motion. Rotational symmetry results in the total angular momentum being constant
 \mathbf{A} = \sum_{i=1}^N \mathbf{q}_i \times \mathbf{p}_i,
where \times is the cross product. The three components of the total angular momentum \mathbf{A} yield three more constants of the motion. The last general constant of the motion is given by the conservation of energy H. Hence, every nbody problem has ten integrals of motion.
Because T and U are homogeneous functions of degree 2 and −1, respectively, the equations of motion have a scaling invariance: if \mathbf{q}_i(t) is a solution, then so is \lambda^{2/3}\mathbf{q}_i(\lambda t) for any \lambda > 0.^{[16]}
The moment of inertia of an nbody system is given by
 I = \sum_{i=1}^N m_i \mathbf{q}_i \cdot \mathbf{q}_i = \sum_{i=1}^N m_i \\mathbf{q}_i\^2
and the virial is given by Q = (1/2)dI/dt. Then the LagrangeJacobi formula states that^{[17]}
 \frac{d^2I}{dt^2} = 2T  U.
For systems in dynamic equilibrium, the longterm time average of \langle d^2I/dt^2 \rangle is zero. Then on average the total kinetic energy is half the total potential energy, \langle T\rangle = \langle U\rangle /2, which is an example of the virial theorem for gravitational systems.^{[18]} If M is the total mass and R a characteristic size of the system (for example, the radius containing half the mass of the system), then the critical time for a system to settle down to a dynamic equilibrium is t_{\rm cr} = \sqrt{GM/R^3}.^{[19]}
Special cases
Twobody problem
Any discussion of planetary interactive forces has always started historically with the Twobody problem. The purpose of this Section is to relate the real complexity in calculating any planetary forces. Note in this Section also, several subjects, such as gravity, barycenter, Kepler's Laws, etc.; and in the following Section too (Threebody problem) — are discussed on other WorldHeritage pages. Here though, these subjects are discussed from the perspective of the nbody problem.
The twobody problem (N=2) was completely solved by Johann Bernoulli (16671748) by classical theory (and not by Newton) by assuming the main pointmass was fixed, is outlined here.^{[20]} Consider then the motion of two bodies, say SunEarth, with the Sun fixed, then:
 m_1 \mathbf{a}_1 = \frac{Gm_1m_2}{r_{12}^3}(\mathbf{r}_2\mathbf{r}_1)\quad\text{SunEarth}
 m_2 \mathbf{a}_2 = \frac{Gm_1m_2}{r_{21}^3}(\mathbf{r}_1\mathbf{r}_2)\quad\text{EarthSun}
The equation describing the motion of mass m_2 relative to mass m_1 is readily obtained from the differences between these two equations and after canceling common terms gives: \alpha + (\eta / r^3)\mathbf{r} = 0, where
 \mathbf{r} = \mathbf{r}_2  \mathbf{r}_1 is the vector position of m_2 relative to m_1;
 \alpha is the Eulerian acceleration d^2\mathbf{r}/dt^2;
 and \eta = G(m_1+m_2).
The equation \alpha + (\eta / r^3)\mathbf{r} = 0 is the fundamental differential equation for the twobody problem Bernoulli solved in 1734. Notice for this approach forces have to be determined first, then the equation of motion resolved. This differential equation has elliptic, or parabolic or hyperbolic solutions^{[21]}^{,} ^{[22]}^{,}.^{[23]}
It is incorrect to think of m_1 (the Sun) as fixed in space when applying Newton's Law of Universal Gravitation, and to do so leads to erroneous results. The fixed point for two isolated gravitationally interacting bodies is their mutual barycenter, and this twobody problem can be solved exactly, such as using Jacobi coordinates relative to the barycenter.
Dr. Clarence Cleminshaw calculated the approximate position of the Solar System's barycenter, a result achieved mainly by combining only the masses of Jupiter and the Sun. Science Program stated in reference to his work:
 "The Sun contains 98 per cent of the mass in the solar system, with the superior planets beyond Mars accounting for most of the rest. On the average, the center of the mass of the SunJupiter system, when the two most massive objects are considered alone, lies 462,000 miles from the Sun's center, or some 30,000 miles above the solar surface! Other large planets also influence the center of mass of the solar system, however. In 1951, for example, the systems' center of mass was not far from the Sun's center because Jupiter was on the opposite side from Saturn, Uranus and Neptune. In the late 1950s, when all four of these planets were on the same side of the Sun, the system's center of mass was more than 330,000 miles from the solar surface, Dr. C. H. Cleminshaw of Griffith Observatory in Los Angeles has calculated."^{[24]}
The Sun wobbles as it rotates around the galactic center, dragging the Solar System and Earth along with it. What mathematician Kepler did in arriving at his three famous equations was curvefit the apparent motions of the planets using Tycho Brahe's data, and not curvefitting their true circular motions about the Sun (see Figure). Both Robert Hooke and Newton were well aware that Newton's Law of Universal Gravitation did not hold for the forces associated with elliptical orbits.^{[11]} In fact, Newton's Universal Law does not account for the orbit of Mercury, the asteroid belt's gravitational behavior, or Saturn's rings.^{[25]} Newton stated (in the 11th Section of the Principia) that the main reason, however, for failing to predict the forces for elliptical orbits was that his math model was for a body confined to a situation that hardly existed in the real world, namely, the motions of bodies attracted toward an unmoving center. Some present physics and astronomy textbooks don't emphasize the negative significance of Newton's assumption and end up teaching that his math model is in effect reality. It is to be understood that the classical twobody problem solution above is a mathematical idealization. See also Kepler's first law of planetary motion.
Some modern writers have criticized Newton's fixed Sun as emblematic of a school of reductive thought  see Truesdell's Essays in the History of Mechanics referenced below. An aside: Newtonian physics doesn't include (among other things) relative motion and may be the root of the reason Newton "fixed" the Sun.^{[26]}^{[27]}
Threebody problem
This Section relates an historically important nbody problem solution after simplifying assumptions were made.
In the past not much was known about the nbody problem for n equal to or greater than three.^{[28]} The case for n = 3 has been the most studied. Many earlier attempts to understand the Threebody problem were quantitative, aiming at finding explicit solutions for special situations.
 In 1687 Isaac Newton published in the Principia the first steps in the study of the problem of the movements of three bodies subject to their mutual gravitational attractions, but his efforts resulted in verbal descriptions and geometrical sketches; see especially Book 1, Proposition 66 and its corollaries (Newton, 1687 and 1999 (transl.), see also Tisserand, 1894).
 In 1767 Euler found collinear motions, in which three bodies of any masses move proportionately along a fixed straight line. The circular restricted threebody problem is the special case in which two of the bodies are in circular orbits (approximated by the SunEarthMoon system and many others).
 In 1772 Lagrange discovered two classes of periodic solution, each for three bodies of any masses. In one class, the bodies lie on a rotating straight line. In the other class, the bodies lie at the vertices of a rotating equilateral triangle. In either case, the paths of the bodies will be conic sections. Those solutions led to the study of central configurations , for which \ddot q=kq for some constant k>0 .
 A major study of the EarthMoonSun system was undertaken by CharlesEugène Delaunay, who published two volumes on the topic, each of 900 pages in length, in 1860 and 1867. Among many other accomplishments, the work already hints at chaos, and clearly demonstrates the problem of socalled "small denominators" in perturbation theory.
 In 1917 Forest Ray Moulton published his now classic, An Introduction to Celestial Mechanics (see references) with its plot of the Restricted Threebody Problem solution (see figure below).^{[29]} An aside, see Meirovitch's book, pages 414 and 413 for his Restricted Threebody Problem solution.^{[30]}
Moulton's solution may be easier to visualize (and definitely easier to solve) if one considers the more massive body (e.g., Sun) to be "stationary" in space, and the less massive body (e.g., Jupiter) to orbit around it, with the equilibrium points (Lagrangian points) maintaining the 60 degreespacing ahead of, and behind, the less massive body almost in its orbit (although in reality neither of the bodies are truly stationary, as they both orbit the center of mass of the whole system—about the barycenter). For sufficiently small mass ratio of the primaries, these triangular equilibrium points are stable, such that (nearly) massless particles will orbit about these points as they orbit around the larger primary (Sun). The five equilibrium points of the circular problem are known as the Lagrangian points. See figure below:
In the Restricted 3Body Problem math model figure above (ref. Moulton), the Lagrangian points L_{4} and L_{5} are where the Trojan planetoids resided (see Lagrangian point); m_{1} is the Sun and m_{2} is Jupiter. L_{2} is where the asteroid belt is. It has to be realized for this model, this whole SunJupiter diagram is rotating about its barycenter. The Restricted 3Body Problem solution predicted the Trojan planetoids before they were first seen. The hcircles and closed loops echo the electromagnetic fluxes issued from the Sun and Jupiter. It is conjectured, contrary to Richard H. Batin conjecture (see References), the two h_{1}'s are gravity sinks, in and where gravitational forces are zero, and the reason the Trojan planetoids are trap there. The total amount of mass of the planetoids is unknown.
The Restricted Threebody Problem assumes the mass of one of the bodies is negligible. For a discussion of the case where the negligible body is a satellite of the body of lesser mass, see Hill sphere; for binary systems, see Roche lobe. Specific solutions to the Threebody problem result in chaotic motion with no obvious sign of a repetitious path.
The restricted problem (both circular and elliptical) was worked on extensively by many famous mathematicians and physicists, most notably by Poincaré at the end of the 19th century. Poincaré's work on the Restricted Threebody Problem was the foundation of deterministic chaos theory. In the restricted problem, there exist five equilibrium points. Three are collinear with the masses (in the rotating frame) and are unstable. The remaining two are located on the third vertex of both equilateral triangles of which the two bodies are the first and second vertices.
Planetary problem
The planetary problem is the nbody problem in the case that one of the masses is much larger than all the others. A prototypical example of a planetary problem is the SunJupiterSaturn system, where the mass of the Sun is about 1000 times bigger than the masses of Jupiter or Saturn.^{[16]} An approximate solution to the problem is to decompose it into N1 pairs of starplanet Kepler problems, treating interactions among the planets as perturbations. Perturbative approximation works well as long as there are no orbital resonances in the system, that is none of the ratios of unperturbed Kepler frequencies is a rational number. Resonances appear as small denominators in the expansion.
The existence of resonances and small denominators lead to the important question of stability in the planetary problem: do planets, in nearly circular orbits around a star, remain in stable or bounded orbits over time?^{[16]}^{[31]} In 1963, Vladimir Arnold proved using KAM theory a kind of stability of the planetary problem: there exists a set of positive measure of quasiperiodic orbits in the case of the planetary problem restricted to the plane.^{[31]} In the KAM theory, chaotic planetary orbits would be bounded by quasiperiodic KAM tori. Arnold's result was extended to a more general theorem by Féjoz and Herman in 2004.^{[32]}
Central configurations
A central configuration \mathbf{q}_1(0), \ldots, \mathbf{q}_N(0) is an initial configuration such that if the particles were all released with zero velocity, they would all collapse toward the center of mass \mathbf{C}.^{[31]} Such a motion is called homothetic. Central configurations may also give rise to homographic motions in which all masses moves along Keplerian trajectories (elliptical, circular, parabolic, or hyperbolic), with all trajectories having the same eccentricity e. For elliptical trajectories, e=1 corresponds to homothetic motion and e=0 gives a relative equilibrium motion in which the configuration remains an isometry of the initial configuration, as if the configuration was a rigid body.^{[33]} Central configurations have played an important role in understanding the topology of invariant manifolds created by fixing the first integrals of a system.
nbody choreography
Solutions in which all masses move on the same curve without collisions are called choreographies.^{[34]} A choreography for N=3 was discovered by Lagrange in 1772 in which three bodies are situated at the vertices of an equilateral triangle in the rotating frame. A figure eight choreography for N=3 was found numerically by C. Moore in 1993 and generalized and proven by A. Chenciner and R. Montgomery in 2000. Since then, many other choreographies have been found for N\ge 3.
Analytic approaches
For every solution of the problem, not only applying an isometry or a time shift but also a reversal of time (unlike in the case of friction) gives a solution as well.
In the physical literature about the nbody problem (n ≥ 3), sometimes reference is made to the impossibility of solving the nbody problem (via employing the above approach). However, care must be taken when discussing the 'impossibility' of a solution, as this refers only to the method of first integrals (compare the theorems by Abel and Galois about the impossibility of solving algebraic equations of degree five or higher by means of formulas only involving roots).
Power series solution
One way of solving the classical nbody problem is "the nbody problem by Taylor series", which is an implementation of the Power series solution of differential equations.
We start by defining the system of differential equations:
As x_{i} (t = t_{0}) and dx_{i}(t)/dt_{t=t0} are given as initial conditions, every d^{2}x_{i}(t)/dt^{2} are known. Differentiating d^{2}x_{i}(t)/dt^{2} results in d^{3}x_{i}(t)/dt^{3} which at t_{0} which is also known, and the Taylor series is constructed iteratively.
A generalized Sundman global solution
In order to generalize Sundman's result for the case n > 3 (or n = 3 and c = 0) one has to face two obstacles:
 As it has been shown by Siegel, collisions which involve more than two bodies cannot be regularized analytically, hence Sundman's regularization cannot be generalized.
 The structure of singularities is more complicated in this case: other types of singularities may occur (see below).
Lastly, Sundman's result was generalized to the case of n > 3 bodies by Q. Wang in the 1990s. Since the structure of singularities is more complicated, Wang had to leave out completely the questions of singularities. The central point of his approach is to transform, in an appropriate manner, the equations to a new system, such that the interval of existence for the solutions of this new system is [0,\infty).
Singularities of the nbody problem
There can be two types of singularities of the nbody problem:
 collisions of two or more bodies, but for which q(t) (the bodies' positions) remains finite. (In this mathematical sense, a "collision" means that two pointlike bodies have identical positions in space.)
 singularities in which a collision does not occur, but q(t) does not remain finite. In this scenario, bodies diverge to infinity in a finite time, while at the same time tending towards zero separation (an imaginary collision occurs "at infinity").
The latter ones are called Painlevé's conjecture (nocollisions singularities). Their existence has been conjectured for n > 3 by Painlevé (see Painlevé conjecture). Examples of this behavior for n=5 have been constructed by Xia^{[35]} and a heuristic model for n=4 by Gerver.^{[36]} Donald G. Saari has shown that singularities have measure zero for many cases.^{[37]}
Simulation
While there are analytic solutions available for the twobody problem and for selected configurations with N > 2, in general nbody problems must be solved or simulated using numerical methods.^{[19]}
Few bodies
For a small number of bodies, an nbody problem can be solved using direct methods, also called particleparticle methods. These methods numerically integrate the differential equations of motion. Numerical integration for this problem can be a challenge for several reasons. First, the gravitational potential is singular; it goes to infinity as the distance between two particles goes to zero. The gravitational potential may be softened to remove the singularity at small distances:^{[19]}
 U_\epsilon = \sum_{1 \le i < j \le N} \frac{G m_i m_j}{\left( \left\ \mathbf{q}_j  \mathbf{q}_i\right\^2 + \epsilon \right)^{1/2}}.
Second, in general for N > 2, the Nbody problem is chaotic,^{[38]} which means that even small errors in integration may grow exponentially in time. Third, a simulation may be over large stretches of model time (e.g., millions of years) and numerical errors accumulate as integration time increases.
There are a number of techniques to reduce errors in numerical integration.^{[19]} Local coordinate systems are used to deal with widely differing scales in some problems, for example an EarthMoon coordinate system in the context of a solar system simulation. Variational methods and perturbation theory can yield approximate analytic trajectories upon which the numerical integration can be a correction. The use of a symplectic integrator ensures that the simulation obeys Hamilton's equations to a high degree of accuracy and in particular that energy is conserved.
Many bodies
Direct methods using numerical integration require on the order of N^2/2 computations to evaluate the potential energy over all pairs of particles, and thus have a time complexity of O(N^2). For simulations with many particles, the O(N^2) factor makes largescale calculations especially time consuming.^{[19]}
A number of approximate methods have been developed that reduce the time complexity relative to direct methods:^{[19]}
 Tree code methods, such as a BarnesHut simulation, are collisionless methods used when close encounters among pairs are not important and distant particle contributions do not need to be computed to high accuracy. The potential of a distant group of particles is computed using a multipole expansion of the potential. This approximation allows for a reduction in complexity to O(N\log(N)).
 Fast multipole methods take advantage of the fact that the multipoleexpanded forces from distant particles are similar for particles close to each other. It is claimed that this further approximation reduces the complexity to O(N).^{[19]}
 Particle mesh methods divide up simulation space into a three dimensional grid onto which the mass density of the particles is interpolated. Then calculating the potential becomes a matter of solving a Poisson equation on the grid, which can be computed in O(N\log(N)) time using fast Fourier transform techniques. Using adaptive mesh refinement or multigrid techniques can further reduce the complexity of the methods.
 P3M and PMTree methods are hybrid methods that use the particle mesh approximation for distant particles, but use more accurate methods for close particles (within a few grid intervals). P3M stands for P^3M or ParticleParticleParticleMesh and uses direct methods with softened potentials at close range. PMTree methods instead use tree codes at close range. As with particle mesh methods, adaptive meshes can increase computational efficiency.
 Mean field methods approximate the system of particles with a timedependent Boltzmann equation representing the mass density that is coupled to a selfconsistent Poisson equation representing the potential. It is a type of smoothedparticle hydrodynamics approximation suitable for large systems.
Strong gravitation
In astrophysical systems with strong gravitational fields, such as those near the event horizon of a black hole, nbody simulations must take into account general relativity; such simulations are the domain of numerical relativity. Numerically simulating the Einstein field equations is extremely challenging^{[19]} and a parameterized postNewtonian formalism (PPN), such as the Einstein–Infeld–Hoffmann equations, is used if possible. The twobody problem in general relativity is analytically solvable only for the Kepler problem, in which one mass is assumed to be much larger than the other.^{[39]}
Other nbody problems
Most work done on the nbody problem has been on the gravitational problem. But there exist other systems for which nbody mathematics and simulation techniques have proven useful.
In large scale electrostatics problems, such as the simulation of proteins and cellular assemblies in structural biology, the Coulomb potential has the same form as the gravitational potential, except that charges may be positive or negative, leading to repulsive as well as attractive forces.^{[40]} Fast Coulomb solvers are the electrostatic counterpart to fast multipole method simulators. These are often used with periodic boundary conditions on the region simulated and Ewald summation techniques are used to speed up computations.^{[41]}
In statistics and machine learning, some models have loss functions of a form similar to that of the gravitational potential: a sum of kernel functions over all pairs of objects, where the kernel function depends on the distance between the objects in parameter space.^{[42]} Example problems that fit into this form include allnearestneighbors in manifold learning, kernel density estimation, and kernel machines. Alternative optimizations to reduce the O(N^2) time complexity to O(N) have been developed, such as dual tree algorithms, that have applicability to the gravitational nbody problem as well.
See also
 Celestial mechanics
 Gravitational twobody problem
 Jacobi integral
 Lunar theory
 Natural units
 Numerical model of the Solar System
 Stability of the Solar System
Notes
 ^ Leimanis and Minorsky: Our interest is with Leimanis, who first discusses some history about the nbody problem, especially Ms. Kovalevskaya's ~18681888, twentyyear complexvariables approach, failure; Section 1: The Dynamics of Rigid Bodies and Mathematical Exterior Ballistics (Chapter 1, the motion of a rigid body about a fixed point (Euler and Poisson equations); Chapter 2, Mathematical Exterior Ballistics), good precursor background to the nbody problem; Section 2: Celestial Mechanics (Chapter 1, The Uniformization of the Threebody Problem (Restricted Threebody Problem); Chapter 2, Capture in the ThreeBody Problem; Chapter 3, Generalized nbody Problem).
 ^ See references cited for Heggie and Hut.
 ^ Quasisteady loads refers to the instantaneous inertial loads generated by instantaneous angular velocities and accelerations, as well as translational accelerations (9 variables). It is as though one took a photograph, which also recorded the instantaneous position and properties of motion. In contrast, a steadystate condition refers to a system's state being invariant to time; otherwise, the first derivatives and all higher derivatives are zero.
 ^ R. M. Rosenberg states the nbody problem similarly (see References): Each particle in a system of a finite number of particles is subjected to a Newtonian gravitational attraction from all the other particles, and to no other forces. If the initial state of the system is given, how will the particles move? Rosenberg failed to realize, like everyone else, that it is necessary to determine the forces first before the motions can be determined.
 ^ A general, classical solution in terms of first integrals is known to be impossible. An exact theoretical solution for arbitrary n can be approximated via Taylor series, but in practice such an infinite series must be truncated, giving at best only an approximate solution; and an approach now obsolete. In addition, the nbody problem may be solved using numerical integration, but these, too, are approximate solutions; and again obsolete. See Sverre J. Aarseth's book Gravitational Nbody Simulations listed in the References.
 ^ See David H. and Stephen P. H. Clark's The Suppressed Scientific Discoveries of Stephen Gray and John Flamsteed, Newton's Tyranny, W. H. Freeman and Co., 2001. A popularization of the historical events and bickering between those parties, but more importantly about the results they produced.
 ^ See "Discovery of gravitation, A.D. 1666" by Sir David Brewster, in The Great Events by Famous Historians, Rossiter Johnson, LL.D. EditorinChief, Volume XII, pp. 5165, The National Alumni, 1905.
 ^ Rudolf Kurth has an extensive discussion in his book (see References) on planetary perturbations.
 ^ An aside: these mathematically undefined planetary perturbations (wobbles) still exist undefined even today and planetary orbits have to be constantly updated, usually yearly. See Astronomical Ephemeris and the American Ephemeris and Nautical Almanac, prepared jointly by the Nautical Almanac Offices of the United Kingdom and the United States of America.
 ^ See Principia, Book Three, System of the World, "General Scholium," page 372, last paragraph. Newton was well aware his math model did not reflect physical reality. This edition referenced is from the Great Books of the Western World, Volume 34, which was translated by Andrew Motte and revised by Florian Cajori. This same paragraph is on page 1160 in Stephen Hawkins' huge On the Shoulders of Giants, 2002 edition; is a copy from Daniel Adee's 1848 addition. Cohen also has translated new editions: Introduction to Newton's 'Principia' , 1970; and Isaac Newton's Principia, with Varian Readings, 1972. Cajori also wrote a History of Science, which is on the Internet.
 ^ ^{a} ^{b} See. I. Bernard Cohen's Scientific American article.
 ^ for details of the serious error in Poincare's first submission see the article by Diacu
 ^ Meyer 2009, pp. 2728
 ^ Meyer 2009, p. 28
 ^ Meyer 2009, pp. 2829
 ^ ^{a} ^{b} ^{c} Chenciner 2007
 ^ Meyer 2009, p. 34
 ^ "AST1100 Lecture Notes: 5 The virial theorem". University of Oslo. Retrieved 25 March 2014.
 ^ ^{a} ^{b} ^{c} ^{d} ^{e} ^{f} ^{g} ^{h} Trenti 2008
 ^ See Bate, Mueller, and White: Chapter 1, "TwoBody Orbital Mechanics," pp 149. These authors were from the Dept. of Astronautics and Computer Science, United States Air Force Academy. See Chapter 1. Their textbook is not filled with advanced mathematics.

^ For the classical approach, if the common center of mass (i.e., the barycenter) of the two bodies is considered to be at rest, then each body travels along a conic section which has a focus at the barycenter of the system. In the case of a hyperbola it has the branch at the side of that focus. The two conics will be in the same plane. The type of conic (circle, ellipse, parabola or hyperbola) is determined by finding the sum of the combined kinetic energy of two bodies and the potential energy when the bodies are far apart. (This potential energy is always a negative value; energy of rotation of the bodies about their axes is not counted here)
 If the sum of the energies is negative, then they both trace out ellipses.
 If the sum of both energies is zero, then they both trace out parabolas. As the distance between the bodies tends to infinity, their relative speed tends to zero.
 If the sum of both energies is positive, then they both trace out hyperbolas. As the distance between the bodies tends to infinity, their relative speed tends to some positive number.
 ^ For this approach see Lindsay's Physical Mechanics, Chapter 3, "Curvilinear Motion in a Plane," and specifically paragraph 39, "Planetary Motion"; and continue reading on to the Chapter's end, pp. 8396. Lindsay presentation goes a long way in explaining these latter comments for the fixed 2body problem; i.e., when the Sun is assumed fixed.
 ^ Note: The fact a parabolic orbit has zero energy arises from the assumption the gravitational potential energy goes to zero as the bodies get infinitely far apart. One could assign any value to the potential energy in the state of infinite separation. That state is assumed to have zero potential energy by convention.
 ^ Science Program's The Nature of the Universe states Clarence Cleminshaw (19021985) served as Assistant Director of Griffith Observatory from 19381958 and as Director from 19581969. Some publications by Cleminshaw, C. H.: Celestial Speeds, 4 1953, equation, Kepler, orbit, comet, Saturn, Mars, velocity; Cleminshaw, C. H.: The Coming Conjunction of Jupiter and Saturn, 7 1960, Saturn, Jupiter, observe, conjunction; Cleminshaw, C. H.: The Scale of The Solar System, 7 1959, Solar system, scale, Jupiter, sun, size, light.
 ^ Brush, Stephen G. Editor: Maxwell on Saturn's Rings, MIT Press, 1983.
 ^ See Jacob Bronowski and Bruce Mazlish's The Western Intellectual Tradition, Dorset Press, 1986, for a discussion of this apparent lack of understanding by Newton. Also see Truesdell's Essays in the History of Mechanics for additional background about Newton accomplishments or lack therein.
 ^ As Hufbauer points out, Newton miscalculated and published unfortunately the wrong value for the Sun's mass twice before he got it correct in his third attempt.
 ^ See Leimanis and Minorsky's historical comments.
 ^ See Moulton's Restricted Threebody Problem 's analytical and graphical solution.
 ^ See Meirovitch's book: Chapters 11, Problems in Celestial Mechanics; 12, Problem in Spacecraft Dynamics; and Appendix A: Dyadics.
 ^ ^{a} ^{b} ^{c} Chierchia 2010
 ^ Féjoz 2004
 ^ See Chierchia 2010 for animations illustrating homographic motions
 ^ Celletti 2008
 ^ Zhihong Xia. The Existence of Noncollision Singularities in Newtonian Systems. Annals of Mathematics. Second Series, Vol. 135, No. 3 (May, 1992), pp. 411468
 ^ Joseph L. Gerver, Noncollision Singularities: Do Four Bodies Suffice?, Exp. Math. (2003), 187198.
 ^ Donald G. Saari. A global existence theorem for the fourbody problem of Newtonian mechanics, J. Differential Equations 26 (1977), 80–111.
 ^ Alligood 1996
 ^ Blanchet 2001
 ^ Krumscheid 2010
 ^ Board 1999
 ^ Ram 2010
References
 Aarseth, Sverre J.: Gravitational Nbody Simulations, Tools and Algorithms, Cambridge University Press, 2003.
 Alligood, K. T., Sauer, T. D., and Yorke, J. A., Chaos: An introduction to Dynamical systems, Springer, pp. 46–48, 1996.
 Bate, Roger R.; Mueller, Donald D.; and White, Jerry: Fundamentals of Astrodynamics, Dover, 1971.
 Blanchet, Luc. On the twobody problem in general relativity, Comptes Rendus de l'Académie des SciencesSeries IVPhysics 2, no. 9 (2001): 13431352.
 Board, John A. Jr (1999), Humphres, Christopher W., Lambert, Christophe G., Rankin, William T., and Toukmaji, Abdulnour Y., Ewald and Multipole Methods for Periodic NBody Problems, in the book Computational Molecular Dynamics: Challenges, Methods, Ideas by editors Deuflhard, Peter, Hermans, Jan, Leimkuhler, Benedict, Mark, Alan E., Reich, Sebastian, and Skeel, Robert D., Springer Berlin Heidelberg, pp. 459–471, http://dx.doi.org/10.1007/9783642583605_27 ISBN 9783540632429.
 Bronowski, Jacob and Mazlish, Bruce: The Western Intellectual Tradition, from Leonardo to Hegel, Dorsey Press. 1986.
 Celletti, Alessandra (2008), Computational celestial mechanics, Scholarpedia, 3(9):4079, doi:10.4249/scholarpedia.4079
 Chenciner, Alain (2007), Three body problem, Scholarpedia, 2(10):2111, doi:10.4249/scholarpedia.2111
 Chierchia, Luigi and Mather, John N. (2010), KolmogorovArnoldMoser Theory, Scholarpedia, 5(9):2123, doi:10.4249/scholarpedia.2123
 Cohen, I. Bernard: "Newton's Discovery of Gravity," Scientific American, pp. 167–179, Vol. 244, No. 3, Mar. 1980.
 Cohen, I. Bernard: The Birth of a New Physics, Revised and Updated, W.W. Norton & Co., 1985.
 Diacu, F.: The solution of the nbody problem, The Mathematical Intelligencer,1996,18,p. 66–70
 Féjoz, J (2004). Dèmonstration du `théorème d'Arnol'd' sur la stabilité du système planétaire (d'après Herman), Ergodic Theory Dynam. Systems, vol 5 : 15211582.
 Heggie, Douglas and Hut, Piet: The Gravitational MillionBody Problem, A Multidisciplinary Approach to Star Cluster Dynamics, Cambridge University Press, 357 pages, 2003.
 Heggie, Douglas C.: "Chaos in the Nbody Problem of Stellar Dynamics," in Predictability, Stability and Chaos in NBody Dynamical Systems, Ed. by Roy A. E., Plenum Press, 1991.
 Hufbauer, Dr. Karl C. (History of Science): Exploring the Sun, Solar Science since Galileo, The Johns Hopkins University Press, 1991. This book was sponsored by the NASA History Office.
 Krumscheid, Sebastian (2010), Benchmark of fast Coulomb Solvers for open and periodic boundary conditions, Jülich Supercomputing Centre, Technical Report FZJJSCIB201001.
 Kurth, Rudolf, Introduction to the Mechanics of the Solar System, Pergamon Press, 1959.
 Leimanis, E., and Minorsky, N.: Dynamics and Nonlinear Mechanics, Part I: Some Recent Advances in the Dynamics of Rigid Bodies and Celestial Mechanics (Leimanis), Part II: The Theory of Oscillations (Minorsky), John Wiley & Sons, Inc., 1958.
 Lindsay, Robert Bruce: Physical Mechanics, 3rd Ed., D. Van Nostrand Co., Inc., 1961.
 Meirovitch, Leonard: Methods of Analytical Dynamics, McGrawHill Book Co., 1970.
 Meyer, Kenneth Ray and Hall, Glen R., Introduction to Hamiltonian Dynamical Systems and the Nbody Problem, Springer Science+Business Media, LLC, 2009, DOI 10.1007/9780387097244 2.
 MittagLeffler, G.: The nbody problem (Price Announcement), Acta Matematica, 1885/1886,7
 Moulton, Forest Ray: An Introduction to Celestial Mechanics, Dover, 1970.
 Newton, I.: Philosophiae Naturalis Principia Mathematica, London, 1687: also English translation of 3rd (1726) edition by I. Bernard Cohen and Anne Whitman (Berkeley, CA, 1999).
 Ram, Parikshit (2009), Dongryeol Lee, William B. March, and Alexander G. Gray. Lineartime Algorithms for Pairwise Statistical Problems, in NIPS, pp. 1527–1535. 2009.
 Rosenberg, Reinhardt M.: Analytical Dynamics, of Discrete Systems, Chapter 19, About Celestial Problems, paragraph 19.5, The nbody Problem, pp. 364–371, Plenum Press, 1977. Like Battin above, Rosenberg employs energy methods too, and to the solution of the general nbody problem but doesn't actually solve anything.
 Science Program's The Nature of the Universe, booklet, published by Nelson Doubleday, Inc., 1968.
 Sundman, K. F.: Memoire sur le probleme de trois corps, Acta Mathematica 36 (1912): 105–179.
 Tisserand, FF.: Mecanique Celeste, tome III (Paris, 1894), ch.III, at p. 27.
 Trenti, Michele and Hut, Piet (2008), Nbody simulations, Scholarpedia, 3(5):3930, doi:10.4249/scholarpedia.3930
 Truesdell, Clifford: Essays in the History of Mechanics, SpringerVerlay, 1968.
 van Winter, Clasine: The nbody problem on a Hilbert space of analytic functions, Paper 1129, in Analytic Methods in Mathematical Physics, edited by Robert P. Gilbert and Roger G. Newton, pp. 569–578, Gordon and Breach, 1970.
 Xia, Zhihong (1992). "The Existence of Noncollision Singularities in Newtonian Systems".
Further reading
 Brouwer, Dirk and Clemence, Gerald M.: Methods of Celestial Mechanics, Academic Press, 1961.
 Battin, Richard H.: An Introduction to The Mathematics and Methods of Astrodynamics, AIAA, 1987. He employs energy methods rather than a Newtonian approach.
 Gelman, Harry: Part I: The second orthogonality conditions in the theory of proper and improper rotations: Derivation of the conditions and of their main consequences, J. Res. NBS 72B (Math. Sci.)No. 3, 1968. Part II: The intrinsic vector; Part III: The Conjugacy Theorem,J. Res. NBS 72B (Math. Sci.) No. 2, 1969. A Note on the time dependence of the effective axis and angle of a rotation, J. Res. NBS 72B (Math. Sci.)No. 3&4, Oct. 1971. These papers are on the Internet.
 Meriam, J. L.: Engineering Mechanics, Volume 1 Statics, Volume 2 Dynamics, John Wiley & Sons, 1978.
 Quadling, Henley: "Gravitational NBody Simulation: 16 bit DOS version," June 1994. nbody*.zip is available at the http://www.ftp.cica.indiana.edu see external links.
 Korenev, G. V.: The Mechanics of Guided Bodies, CRC Press, 1967.
 Eisele, John A. and Mason, Robert M.: Applied Matrix and Tensor Analysis, John Wiley & Sons, 1970.
 Murray, Carl D. and Dermott, Stanley F.: Solar System Dynamics, Cambridge University Press, 606 pages, 2000.
 Crandall, Richard E.: Topics in Advanced Scientific Computation, Chapter 5, "Nonlinear & Complex Systems," paragraph 5.1, "Nbody problems & chaos," pp. 215–221, SpringerVerlag, 1996.
 Crandall, Richard E.: Projects in Scientific Computation, Chapter 2, "Exploratory Computation," Project 2.4.1, "Classical Physics," pp. 93–97, corrected 3rd printing, SpringerVerlag, 1996.
 Szebehely, Victor: Theory of Orbits, Academic Press, 1967.
 Saari, D.: A visit to the Newtonian nbody problem via Elementary Complex Variables, American Mathematical Monthly, 1990, 89, 105–119
 Saari, D. G.; Hulkower, N. D. (1981). "On the Manifolds of Total Collapse Orbits and of Completely Parabolic Orbits for the nBody Problem". Journal of Differential Equations 41 (1): 27–43.
 Hagihara, Y: Celestial Mechanics. (Vol I and Vol II pt 1 and Vol II pt 2.) MIT Press, 1970.
 Boccaletti, D. and Pucacco, G.: Theory of Orbits (two volumes). SpringerVerlag, 1998.
 Havel, Karel. NBody Gravitational Problem: Unrestricted Solution (ISBN 9780968912058). Brampton: Grevyt Press, 2008. http://www.grevytpress.com
External links
 ThreeBody Problem at Scholarpedia
 More detailed information on the threebody problem
 Regular Keplerian motions in classical manybody systems
 Applet demonstrating chaos in restricted threebody problem
 Applets demonstrating many different threebody motions
 body equationsnOn the integration of the
 Java applet simulating Solar System
 Java applet simulating a ring of bodies orbiting a large central mass
 Java applet simulating dust in the Solar System
 Java applet simulating a stable solution to the equimass 3body problem
 Java applet simulating choreographies and other interesting nbody solutions
 A java applet to simulate the 3d movement of set of particles under gravitational interaction
 Javascript Simulation of our Solar System
 The Lagrange Points  with links to the original papers of Euler and Lagrange, and to translations, with discussion
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