Kepler's laws of planetary motion

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This article uses more modern mathematical notation and concepts than were available to Johannes Kepler. For a more historical approach, see in particular the article Astronomia nova.

In astronomy, Kepler's laws of planetary motion are three scientific laws describing the motion of planets around the Sun. Kepler's laws are now traditionally enumerated in this way:

Figure 1: Illustration of Kepler's three laws with two planetary orbits.
(1) The orbits are ellipses, with focal points ƒ1 and ƒ2 for the first planet and ƒ1 and ƒ3 for the second planet. The Sun is placed in focal point ƒ1.

(2) The two shaded sectors A1 and A2 have the same surface area and the time for planet 1 to cover segment A1 is equal to the time to cover segment A2.

(3) The total orbit times for planet 1 and planet 2 have a ratio a13/2 : a23/2.
  1. The orbit of a planet is an ellipse with the Sun at one of the two foci.
  2. A line segment joining a planet and the Sun sweeps out equal areas during equal intervals of time.[1]
  3. The square of the orbital period of a planet is proportional to the cube of the semi-major axis of its orbit.

Most planetary orbits are almost circles, so it is not apparent that they are actually ellipses. Calculations of the orbit of the planet Mars first indicated to Kepler its elliptical shape, and he inferred that other heavenly bodies, including those farther away from the Sun, have elliptical orbits also. Kepler's work broadly followed the heliocentric theory of Nicolaus Copernicus by asserting that the Earth orbited the Sun. It innovated in explaining how the planets' speeds varied, and using elliptical orbits rather than circular orbits with epicycles.[2]

Isaac Newton showed in 1687 that relationships like Kepler's would apply in the solar system to a good approximation, as consequences of his own laws of motion and law of universal gravitation. Together with Newton's theories, Kepler's laws became part of the foundation of modern astronomy and physics.[3]

Nomenclature[edit]

It took nearly two centuries for the current formulation of Kepler's work to take on its settled form. Voltaire's Eléments de la philosophie de Newton (Elements of Newton's Philosophy) of 1738 was the first publication to use the terminology of "laws".[4] The Biographical Encyclopedia of Astronomers in its article on Kepler (p. 620) states that the terminology of scientific laws for these discoveries was current at least from the time of Joseph de Lalande. It was the exposition of Robert Small, in An account of the astronomical discoveries of Kepler (1804) that made up the set of three laws, by adding in the third. Small also claimed, against the history, that these were empirical laws, based on inductive reasoning.[4]

Further, the current usage of "Kepler's Second Law" is something of a misnomer. Kepler had two versions of it, related in a qualitative sense, the "distance law" and the "area law". The "area law" is what became the Second Law in the set of three; but Kepler did himself not privilege it in that way.[5]

History[edit]

Johannes Kepler published his first two laws about planetary motion in 1609, having found them by analyzing the astronomical observations of Tycho Brahe.[2] Kepler's third law was published in 1619.[2]

Kepler in 1622 and Godefroy Wendelin in 1643 noted that Kepler's third law applies to the four brightest moons of Jupiter.[Nb 1] The second law ("area law" form) was contested by Nicolaus Mercator in a book from 1664; but by 1670 he was publishing in its favour in Philosophical Transactions, and as the century proceeded it became more widely accepted.[6] The reception in Germany changed noticeably between 1688, the year in which Newton's Principia was published and was taken to be basically Copernican, and 1690, by which time work of Gottfried Leibniz on Kepler had been published.[7]

First Law[edit]

"The orbit of every planet is an ellipse with the Sun at one of the two foci."
Figure 2: Kepler's first law placing the Sun at the focus of an elliptical orbit
Figure 4: Heliocentric coordinate system (r, θ) for ellipse. Also shown are: semi-major axis a, semi-minor axis b and semi-latus rectum p; center of ellipse and its two foci marked by large dots. For θ = 0°, r = rmin and for θ = 180°, r = rmax.

Mathematically, an ellipse can be represented by the formula:

r=\frac{p}{1+\varepsilon\, \cos\theta},

where (rθ) are polar coordinates, p is the semi-latus rectum, and ε is the eccentricity of the ellipse.

Note that 0 < ε < 1 for an ellipse; in the limiting case ε = 0, the orbit is a circle with the sun at the centre (see section Zero eccentricity below).

For a planet r is the distance from the Sun to the planet, and θ is the angle to the planet's current position from its closest approach, as seen from the Sun.

At θ = 0°, perihelion, the distance is minimum

r_\mathrm{min}=\frac{p}{1+\varepsilon}.

At θ = 90° and at θ = 270°, the distance is \, p.

At θ = 180°, aphelion, the distance is maximum

r_\mathrm{max}=\frac{p}{1-\varepsilon}.

The semi-major axis a is the arithmetic mean between rmin and rmax:

\,r_\max - a=a-r_\min
a=\frac{p}{1-\varepsilon^2}.

The semi-minor axis b is the geometric mean between rmin and rmax:

\frac{r_\max} b =\frac b{r_\min}
b=\frac p{\sqrt{1-\varepsilon^2}}.

The semi-latus rectum p is the harmonic mean between rmin and rmax:

\frac{1}{r_\min}-\frac{1}{p}=\frac{1}{p}-\frac{1}{r_\max}
pa=r_\max r_\min=b^2\,.

The eccentricity ε is the coefficient of variation between rmin and rmax:

\varepsilon=\frac{r_\mathrm{max}-r_\mathrm{min}}{r_\mathrm{max}+r_\mathrm{min}}.

The area of the ellipse is

A=\pi a b\,.

The special case of a circle is ε = 0, resulting in r = p = rmin = rmax = a = b and A = π r2.

Practical examples
Semi-major axis Orbital eccentricity Interfocal distance Location of f2
Jupiter 7,783 mn km 0.048386 75 mn km Roughly half-way between Mercury's (58 mn km) and Venus's (108 mn km) orbital distances.
Saturn 14,267 mn km 0.053862 154 mn km Roughly at Earth orbit distance (150 mn km).

Second law[edit]

"A line joining a planet and the Sun sweeps out equal areas during equal intervals of time."[1]
The same blue area is swept out in a given time. The green arrow is velocity. The purple arrow directed towards the Sun is the acceleration. The other two purple arrows are acceleration components parallel and perpendicular to the velocity

In a small time dt\, the planet sweeps out a small triangle (or, more precisely, a sector) having base line r\, and height r d\theta\, and area dA=\tfrac 1 2\cdot r\cdot r d\theta and so the constant areal velocity is
\frac{dA}{dt}=\tfrac{1}{2}r^2 \frac{d\theta}{dt}. The planet moves faster when it is closer to the Sun.

The area enclosed by the elliptical orbit is \pi ab.\, So the period P\, satisfies

P\cdot \tfrac 12r^2 \frac{d\theta}{dt}=\pi a b

and the mean motion of the planet around the Sun n = {2\pi}/P satisfies

r^2{d\theta} = a b n dt .
Practical examples
Perihelion distance Perihelion speed Aphelion distance Aphelion speed Sector area swept
in one second
Mercury 46.0 mn km 59,000 m/s 69.8 mn km 38,900 m/s 1.3565×1015 m2
for both perihelion and aphelion
Sedna 11,423 bn km 4,640 m/s ~ 140,000 bn km ~ 430 m/s 3.01×1016 m2 (based on perihelion measurements)
2.65×1016 m2 (based on aphelion measurements)
Note: the point speeds are worked out based on v=√GM(2/d−1/a); the "area swept" figure is less accurate for Sedna due to the imprecise degree to which its orbital parameters are currently known.

Third law[edit]

The square of the orbital period of a planet is directly proportional to the cube of the semi-major axis of its orbit.

The third law, published by Kepler in 1619 [1] captures the relationship between the distance of planets from the Sun, and their orbital periods.

Kepler enunciated this third law in a laborious attempt to determine what he viewed as the "music of the spheres" according to precise laws, and express it in terms of musical notation.[8] So it used to be known as the harmonic law.[9]

Mathematically, the law says that the expression  P^2/a^3 has the same value for all the planets in the solar system.

Examples
Period & Semi-major axis
(in relation to Earth)
Period 2 : Semi-major axis 3
(in relation to Earth)
Period & Semi-major axis
(in real terms)
Period 2 : Semi-major axis 3
(in real terms)
Ratio
Venus 0.615 yr, 0.723 AU 0.3785 : 0.3785 19.4 mn secs, 1.082×1011 m 3.769×1014 : 1.267×1033 1 : 3.362×1018
Neptune 164.8 yrs, 30.1 AU 27160 : 27280 5.2 bn secs, 4.498×1012 m 2.704×1019 : 9.103×1037 1 : 3.366×1018

The modern formulation with the constant evaluated reads:

\frac{T^2}{r^3} = \frac{4 \pi^2}{GM}

with

  • T the orbital period of the star
  • M the mass of the star,
  • G the universal gravitational constant and
  • r the radius, the semi-major axis of the ellipse.

In the full formulation under Newton's laws of motion, M should be replaced by (M+m), where m is the mass of the orbiting body. Consequently, the proportionality constant is not truly the same for each planet. Nevertheless, m \ll M for all planets in our solar system such that variations in the proportionality constant are negligible.

Zero eccentricity[edit]

Kepler's laws refine the model of Copernicus, which assumed circular orbits. If the eccentricity of a planetary orbit is zero, then Kepler's laws state:

  1. The planetary orbit is a circle with the Sun at the center
  2. The speed of the planet in the orbit is constant
  3. The square of the sidereal period is proportionate to the cube of the distance from the Sun.

Actually, the eccentricities of the orbits of the six planets known to Copernicus and Kepler are quite small, so the rules above give excellent approximations of planetary motion, but Kepler's laws fit observations even better.

Kepler's corrections to the Copernican model are not at all obvious:

  1. The planetary orbit is not a circle, but an ellipse
  2. The Sun is not at the center but at a focal point
  3. Neither the linear speed nor the angular speed of the planet in the orbit is constant, but the area speed is constant.
  4. The square of the sidereal period is proportionate to the cube of the mean between the maximum and minimum distances from the Sun.

The nonzero eccentricity of the orbit of the earth makes the time from the March equinox to the September equinox, around 186 days, unequal to the time from the September equinox to the March equinox, around 179 days. A diameter would cut the orbit into equal parts, but the plane through the sun parallel to the equator of the earth cuts the orbit into two parts with areas in a 186 to 179 ratio, so the eccentricity of the orbit of the Earth is approximately

\varepsilon\approx\frac \pi 4 \frac {186-179}{186+179}\approx 0.015,

which is close to the correct value (0.016710219). (See Earth's orbit). The calculation is correct when the perihelion, the date that the Earth is closest to the Sun, is on a solstice. The current perihelion, near January 4, is fairly close to the solstice on December 21 or 22.

Planetary acceleration[edit]

A sudden sunward velocity change is applied to a planet. Then the areas of the triangles defined by the path of the planet for fixed time intervals will be equal. (Click on image for a detailed description.)

Isaac Newton computed in his Philosophiæ Naturalis Principia Mathematica the acceleration of a planet moving according to Kepler's first and second law.

  1. The direction of the acceleration is towards the Sun.
  2. The magnitude of the acceleration is in inverse proportion to the square of the distance from the Sun.

This suggests that the Sun may be the physical cause of the acceleration of planets.

Newton defined the force on a planet to be the product of its mass and the acceleration. (See Newton's laws of motion). So:

  1. Every planet is attracted towards the Sun.
  2. The force on a planet is in direct proportion to the mass of the planet and in inverse proportion to the square of the distance from the Sun.

Here the Sun plays an unsymmetrical part, which is unjustified. So he assumed Newton's law of universal gravitation:

  1. All bodies in the solar system attract one another.
  2. The force between two bodies is in direct proportion to the product of their masses and in inverse proportion to the square of the distance between them.

As the planets have small masses compared to that of the Sun, the orbits conform to Kepler's laws approximately. Newton's model improves upon Kepler's model and fits actual observations more accurately. (See two-body problem).

A deviation in the motion of a planet from Kepler's laws due to the gravity of other planets is called a perturbation.

Below comes the detailed calculation of the acceleration of a planet moving according to Kepler's first and second laws.

Acceleration vector[edit]

From the heliocentric point of view consider the vector to the planet \mathbf{r} = r \hat{\mathbf{r}} where  r is the distance to the planet and the direction  \hat {\mathbf{r}} is a unit vector. When the planet moves the direction vector  \hat {\mathbf{r}} changes:

 \frac{d\hat{\mathbf{r}}}{dt}=\dot{\hat{\mathbf{r}}} = \dot\theta  \hat{\boldsymbol\theta},\qquad \dot{\hat{\boldsymbol\theta}} = -\dot\theta \hat{\mathbf{r}}

where \scriptstyle  \hat{\boldsymbol\theta} is the unit vector orthogonal to \scriptstyle \hat{\mathbf{r}} and pointing in the direction of rotation, and \scriptstyle \theta is the polar angle, and where a dot on top of the variable signifies differentiation with respect to time.

So differentiating the position vector twice to obtain the velocity and the acceleration vectors:

\dot{\mathbf{r}} =\dot{r} \hat{\mathbf{r}} + r \dot{\hat{\mathbf{r}}}
=\dot{r} \hat{\mathbf{r}} + r \dot{\theta} \hat{\boldsymbol{\theta}},
\ddot{\mathbf{r}} 
= (\ddot{r} \hat{\mathbf{r}} +\dot{r} \dot{\hat{\mathbf{r}}} )
+ (\dot{r}\dot{\theta} \hat{\boldsymbol{\theta}} + r\ddot{\theta} \hat{\boldsymbol{\theta}}
+ r\dot{\theta} \dot{\hat{\boldsymbol{\theta}}})
= (\ddot{r} - r\dot{\theta}^2) \hat{\mathbf{r}} + (r\ddot{\theta} + 2\dot{r} \dot{\theta}) \hat{\boldsymbol{\theta}}.

So

\ddot{\mathbf{r}} = a_r \hat{\boldsymbol{r}}+a_\theta\hat{\boldsymbol{\theta}}

where the radial acceleration is

a_r=\ddot{r} - r\dot{\theta}^2

and the transversal acceleration is

a_\theta=r\ddot{\theta} + 2\dot{r} \dot{\theta}.

The inverse square law[edit]

Kepler's laws say that

r^2\dot \theta = nab

is constant.

The transversal acceleration a_\theta is zero:

\frac{d (r^2 \dot \theta)}{dt} = r (2 \dot r \dot \theta + r \ddot \theta ) = r a_\theta = 0.

So the acceleration of a planet obeying Kepler's laws is directed towards the sun.

The radial acceleration a_r  is

a_r = \ddot r - r \dot \theta^2= \ddot r - r \left(\frac{nab}{r^2}
\right)^2= \ddot r -\frac{n^2a^2b^2}{r^3}.

Kepler's first law states that the orbit is described by the equation:

\frac{p}{r} = 1+ \varepsilon \cos\theta.

Differentiating with respect to time

-\frac{p\dot r}{r^2} = -\varepsilon  \sin \theta \,\dot \theta

or

p\dot r = nab\,\varepsilon\sin \theta.

Differentiating once more

p\ddot r =nab \varepsilon \cos \theta \,\dot \theta
=nab \varepsilon \cos \theta \,\frac{nab}{r^2}
=\frac{n^2a^2b^2}{r^2}\varepsilon \cos \theta .

The radial acceleration a_r  satisfies

p a_r = \frac{n^2 a^2b^2}{r^2}\varepsilon \cos \theta  - p\frac{n^2 a^2b^2}{r^3}
= \frac{n^2a^2b^2}{r^2}\left(\varepsilon \cos \theta - \frac{p}{r}\right).

Substituting the equation of the ellipse gives

p a_r = \frac{n^2a^2b^2}{r^2}\left(\frac p r - 1 - \frac p r\right)= -\frac{n^2a^2}{r^2}b^2.

The relation b^2=pa gives the simple final result

a_r=-\frac{n^2a^3}{r^2}.

This means that the acceleration vector \mathbf{\ddot r} of any planet obeying Kepler's first and second law satisfies the inverse square law

\mathbf{\ddot r} = - \frac{\alpha}{r^2}\hat{\mathbf{r}}

where

\alpha = n^2 a^3\,

is a constant, and \hat{\mathbf r} is the unit vector pointing from the Sun towards the planet, and r\, is the distance between the planet and the Sun.

According to Kepler's third law, \alpha has the same value for all the planets. So the inverse square law for planetary accelerations applies throughout the entire solar system.

The inverse square law is a differential equation. The solutions to this differential equation include the Keplerian motions, as shown, but they also include motions where the orbit is a hyperbola or parabola or a straight line. See Kepler orbit.

Newton's law of gravitation[edit]

By Newton's second law, the gravitational force that acts on the planet is:

\mathbf{F} = m_{Planet} \mathbf{\ddot r} = - {m_{Planet} \alpha}{r^{-2}}\hat{\mathbf{r}}

where m_{Planet} is the mass of the planet and \alpha has the same value for all planets in the solar system. According to Newton's third Law, the Sun is attracted to the planet by a force of the same magnitude. Since the force is proportional to the mass of the planet, under the symmetric consideration, it should also be proportional to the mass of the Sun, m_{Sun}. So

\alpha = Gm_{Sun}

where G is the gravitational constant.

The acceleration of solar system body number i is, according to Newton's laws:

\mathbf{\ddot r_i} = G\sum_{j\ne i} {m_j}{r_{ij}^{-2}}\hat{\mathbf{r}}_{ij}

where m_j is the mass of body j, r_{ij} is the distance between body i and body j, \hat{\mathbf{r}}_{ij} is the unit vector from body i towards body j, and the vector summation is over all bodies in the world, besides i itself.

In the special case where there are only two bodies in the world, Earth and Sun, the acceleration becomes

\mathbf{\ddot r}_{Earth} = G{m_{Sun}}{r_{{Earth},{Sun}}^{-2}}\hat{\mathbf{r}}_{{Earth},{Sun}}

which is the acceleration of the Kepler motion. So this Earth moves around the Sun according to Kepler's laws.

If the two bodies in the world are Moon and Earth the acceleration of the Moon becomes

\mathbf{\ddot r}_{Moon} = G{m_{Earth}}{r_{{Moon},{Earth}}^{-2}}\hat{\mathbf{r}}_{{Moon},{Earth}}

So in this approximation the Moon moves around the Earth according to Kepler's laws.

In the three-body case the accelerations are

\mathbf{\ddot r}_{Sun} = G m_{Earth}{r_{{Sun},{Earth}}^{-2}}\hat{\mathbf{r}}_{{Sun},{Earth}} + G{m_{Moon}}{r_{{Sun},{Moon}}^{-2}}\hat{\mathbf{r}}_{{Sun},{Moon}}
\mathbf{\ddot r}_{Earth} = G{m_{Sun}}{r_{{Earth},{Sun}}^{-2}}\hat{\mathbf{r}}_{{Earth},{Sun}} + G{m_{Moon}}{r_{{Earth},{Moon}}^{-2}}\hat{\mathbf{r}}_{{Earth},{Moon}}
\mathbf{\ddot r}_{Moon} = G{m_{Sun}}{r_{{Moon},{Sun}}^{-2}}\hat{\mathbf{r}}_{{Moon},{Sun}}+G{m_{Earth}}{r_{{Moon},{Earth}}^{-2}}\hat{\mathbf{r}}_{{Moon},{Earth}}

These accelerations are not those of Kepler orbits, and the three-body problem is complicated. But Keplerian approximation is basis for perturbation calculations. See Lunar theory.

Position as a function of time [edit]

Kepler used his two first laws to compute the position of a planet as a function of time. His method involves the solution of a transcendental equation called Kepler's equation.

The procedure for calculating the heliocentric polar coordinates (r,θ) of a planet as a function of the time t since perihelion, is the following four steps:

1. Compute the mean anomaly M=n\cdot t where n is the mean motion.
n\cdot P=2\pi where P is the period.
2. Compute the eccentric anomaly E by solving Kepler's equation:
\ M=E-\varepsilon\cdot\sin E
3. Compute the true anomaly θ by the equation:
(1-\varepsilon)\cdot\tan^2\frac \theta 2 = (1+\varepsilon)\cdot\tan^2\frac E 2
4. Compute the heliocentric distance r from the first law:
r\cdot(1+\varepsilon\cdot\cos\theta)=a\cdot(1-\varepsilon^2)

The important special case of circular orbit, ε = 0, gives θ = E = M. Because the uniform circular motion was considered to be normal, a deviation from this motion was considered an anomaly.

The proof of this procedure is shown below.

Mean anomaly, M[edit]

FIgure 5: Geometric construction for Kepler's calculation of θ. The Sun (located at the focus) is labeled S and the planet P. The auxiliary circle is an aid to calculation. Line xd is perpendicular to the base and through the planet P. The shaded sectors are arranged to have equal areas by positioning of point y.

The Keplerian problem assumes an elliptical orbit and the four points:

s the Sun (at one focus of ellipse);
z the perihelion
c the center of the ellipse
p the planet

and

\ a=|cz|, distance between center and perihelion, the semimajor axis,
\ \varepsilon={|cs|\over a}, the eccentricity,
\ b=a\sqrt{1-\varepsilon^2}, the semiminor axis,
\ r=|sp| , the distance between Sun and planet.
\theta=\angle zsp, the direction to the planet as seen from the Sun, the true anomaly.

The problem is to compute the polar coordinates (r,θ) of the planet from the time since perihelion, t.

It is solved in steps. Kepler considered the circle with the major axis as a diameter, and

\ x, the projection of the planet to the auxiliary circle
\ y, the point on the circle such that the sector areas |zcy| and |zsx| are equal,
M=\angle zcy, the mean anomaly.

The sector areas are related by |zsp|=\frac b a \cdot|zsx|.

The circular sector area \ |zcy| =  \frac{a^2 M}2.

The area swept since perihelion,

|zsp|=\frac b a \cdot|zsx|=\frac b a \cdot|zcy|=\frac b a\cdot\frac{a^2 M}2 = \frac {a b M}{2},

is by Kepler's second law proportional to time since perihelion. So the mean anomaly, M, is proportional to time since perihelion, t.

M=n t,

where n is the mean motion.

Eccentric anomaly, E[edit]

When the mean anomaly M is computed, the goal is to compute the true anomaly θ. The function θ=f(M) is, however, not elementary.[10] Kepler's solution is to use

E=\angle zcx, x as seen from the centre, the eccentric anomaly

as an intermediate variable, and first compute E as a function of M by solving Kepler's equation below, and then compute the true anomaly θ from the eccentric anomaly E. Here are the details.

\ |zcy|=|zsx|=|zcx|-|scx|
\frac{a^2 M}2=\frac{a^2 E}2-\frac {a\varepsilon\cdot a\sin E}2

Division by a2/2 gives Kepler's equation

M=E-\varepsilon\cdot\sin E.

This equation gives M as a function of E. Determining E for a given M is the inverse problem. Iterative numerical algorithms are commonly used.

Having computed the eccentric anomaly E, the next step is to calculate the true anomaly θ.

True anomaly, θ[edit]

Note from the figure that

\overrightarrow{cd}=\overrightarrow{cs}+\overrightarrow{sd}

so that

a\cdot\cos E=a\cdot\varepsilon+r\cdot\cos \theta.

Dividing by a and inserting from Kepler's first law

\ \frac r a =\frac{1-\varepsilon^2}{1+\varepsilon\cdot\cos \theta}

to get

\cos E
=\varepsilon+\frac{1-\varepsilon^2}{1+\varepsilon\cdot\cos \theta}\cdot\cos \theta
=\frac{\varepsilon\cdot(1+\varepsilon\cdot\cos \theta)+(1-\varepsilon^2)\cdot\cos \theta}{1+\varepsilon\cdot\cos \theta}
=\frac{\varepsilon +\cos \theta}{1+\varepsilon\cdot\cos \theta}.

The result is a usable relationship between the eccentric anomaly E and the true anomaly θ.

A computationally more convenient form follows by substituting into the trigonometric identity:

\tan^2\frac{x}{2}=\frac{1-\cos x}{1+\cos x}.

Get

\tan^2\frac{E}{2}
=\frac{1-\cos E}{1+\cos E}
=\frac{1-\frac{\varepsilon+\cos \theta}{1+\varepsilon\cdot\cos \theta}}{1+\frac{\varepsilon+\cos \theta}{1+\varepsilon\cdot\cos \theta}}
=\frac{(1+\varepsilon\cdot\cos \theta)-(\varepsilon+\cos \theta)}{(1+\varepsilon\cdot\cos \theta)+(\varepsilon+\cos \theta)}
=\frac{1-\varepsilon}{1+\varepsilon}\cdot\frac{1-\cos \theta}{1+\cos \theta}=\frac{1-\varepsilon}{1+\varepsilon}\cdot\tan^2\frac{\theta}{2}.

Multiplying by 1+ε gives the result

(1-\varepsilon)\cdot\tan^2\frac \theta 2 = (1+\varepsilon)\cdot\tan^2\frac E 2

This is the third step in the connection between time and position in the orbit.

Distance, r[edit]

The fourth step is to compute the heliocentric distance r from the true anomaly θ by Kepler's first law:

\ r\cdot(1+\varepsilon\cdot\cos \theta)=a\cdot(1-\varepsilon^2).

See also[edit]

Notes[edit]

  1. ^ Godefroy Wendelin wrote a letter to Giovanni Battista Riccioli about the relationship between the distances of the Jovian moons from Jupiter and the periods of their orbits, showing that the periods and distances conformed to Kepler's third law. See: Joanne Baptista Riccioli, Almagestum novum … (Bologna (Bononiae), (Italy): Victor Benati, 1651), volume 1, page 492 Scholia III. In the margin beside the relevant paragraph is printed: Vendelini ingeniosa speculatio circa motus & intervalla satellitum Jovis. (The clever Wendelin's speculation about the movement and distances of Jupiter's satellites.)
    In 1622, Johannes Kepler had noted that Jupiter's moons obey (approximately) his third law in his Epitome Astronomiae Copernicanae [Epitome of Copernican Astronomy] (Linz (“Lentiis ad Danubium“), (Austria): Johann Planck, 1622), book 4, part 2, page 554.

References[edit]

  1. ^ a b Bryant, Jeff; Pavlyk, Oleksandr. "Kepler's Second Law", Wolfram Demonstrations Project. Retrieved December 27, 2009.
  2. ^ a b c Holton, Gerald James; Brush, Stephen G. (2001). Physics, the Human Adventure: From Copernicus to Einstein and Beyond (3rd paperback ed.). Piscataway, NJ: Rutgers University Press. pp. 40–41. ISBN 0-8135-2908-5. Retrieved December 27, 2009. 
  3. ^ See also G. E. Smith, "Newton's Philosophiae Naturalis Principia Mathematica", especially the section Historical context ... in The Stanford Encyclopedia of Philosophy (Winter 2008 Edition), Edward N. Zalta (ed.).
  4. ^ a b Wilson, Curtis (May 1994). "Kepler's Laws, So-Called". HAD News (Washington, DC: Historical Astronomy Division, American Astronomical Society) (31): 1–2. Retrieved December 27, 2009. 
  5. ^ Bruce Stephenson (1994). Kepler's Physical Astronomy. Princeton University Press. p. 170. ISBN 0-691-03652-7. 
  6. ^ Wilbur Applebaum (13 June 2000). Encyclopedia of the Scientific Revolution: From Copernicus to Newton. Routledge. p. 603. ISBN 978-1-135-58255-5. 
  7. ^ Roy Porter (25 September 1992). The Scientific Revolution in National Context. Cambridge University Press. p. 102. ISBN 978-0-521-39699-8. 
  8. ^ Burtt, Edwin. The Metaphysical Foundations of Modern Physical Science. p. 52.
  9. ^ Gerald James Holton, Stephen G. Brush (2001). Physics, the Human Adventure. Rutgers University Press. p. 45. ISBN 0-8135-2908-5. 
  10. ^ MÜLLER, M (1995). "EQUATION OF TIME -- PROBLEM IN ASTRONOMY". Acta Physica Polonica A. Retrieved 23 February 2013. 

Bibliography[edit]

  • A derivation of Kepler's third law of planetary motion is a standard topic in engineering mechanics classes. See, for example, pages 161–164 of Meriam, J. L. (1966, 1971). Dynamics, 2nd ed. New York: John Wiley. ISBN 0-471-59601-9.  .
  • Murray and Dermott, Solar System Dynamics, Cambridge University Press 1999, ISBN 0-521-57597-4
  • V.I. Arnold, Mathematical Methods of Classical Mechanics, Chapter 2. Springer 1989, ISBN 0-387-96890-3

External links[edit]