Introduction to potential theory – at black board Potentials of simple spherical systems

of 40 /40
theory – at black board mple spherical systems n potential ρ = constant and M(r)=(4/3)πr 3 ρ for r < a for r > a 3

Embed Size (px)


Introduction to potential theory – at black board Potentials of simple spherical systems Point mass - keplerian potential Homogeneous sphere  ρ = constant and M(r)=(4/3)πr 3 ρ With radial size a for r < a for r > a then. 3. - PowerPoint PPT Presentation

Transcript of Introduction to potential theory – at black board Potentials of simple spherical systems

  • Introduction to potential theory at black board

    Potentials of simple spherical systems

    Point mass- keplerian potential

    Homogeneous sphere = constant and M(r)=(4/3)r3With radial size a

    for r < a

    for r > a3

  • Isochrone potential model a galaxy as a constant density at the center with density decreasing at larger radii. One potential with these properties:

    where b is characteristic radius that defines how the density decreases with rDensity pair given in BT (2-34) and yields

    at center and at r >>b Modified Hubble profile derived from SBs for ellipticals

    where a is core radius and j is luminosity density

  • Power-law density profile many galaxies have surface brightness profiles that approximate a power-law over large radii


    we can compute M(r) and Vc(r)

    If = 2, this is an isothermal sphere (density goes as 1/r2)

    Can be used to approximate galaxies with flat rotation curves; need outer cut-off to obtain finite mass

  • Plummer Sphere simple model for round galaxies/clusters

    This potential softens force between particles in N-body simulations by avoiding the singularity of the Newtonian potential. The density profile has finite core density but falls as r-5 at large r (too steep for most galaxies).

    Jaffe and Hernquist profiles

    Both decline as r-4 at large radii which works well with galaxy models produced from violent relaxation (i.e. stellar systems relax quickly from initial state to quasi-equilibrium).

    Hernquist has gentle power-law cusp at small r while Jaffe has steeper cusp. Potential density

  • Density distributions for various simple spherical potentials

  • Navarro, Frenk and White (NFW) profile

    Good fit to dark matter haloes formed in simulationsProblem mass diverges logarithmically with r must be cut off at large rPotentials for Flattened Models: Axisymmetric potential

    Kuzmin Disk (cylindrical coordinates)At points with z0, k is the same as the potential generated by a point mass at (0,-a).Everywhere except on plane z=0

  • Miyamoto & Nagai (1975) introduced a combination Plummer sphere/Kuzmin disk model

    where b is aP in previous Plummer notation

    a=0 Plummer sphereb=0 Kuzmin disk

    b/a ~ 0.2 similar to disk galaxiesUse divergence theorem to find the surface density generated by Kuzmin potential

    Kuzmin (1956) or Toomre model 1 (1962)

  • Stellar Orbits

    For a star moving through a galaxy, assume its motion does not change the overall potentialIf the galaxy is not collapsing, colliding, etc., assume potential does not change with time

    Then, as a star moves with velocity v, the potential at its location changes as

    Recall(grad of potential is force on star)


    Energy along orbit remains constant (KE always + ; PE goes to 0 at large x)

    Star escapes galaxy if E > 0

    Circular velocity angular velocity

  • In a cluster of stars, motions of the stars can cause the potential to change with time. The energy of each individual star is no longer conserved, only the total for the cluster as a whole.

    cluster KEcluster PE

    Stars in a cluster can change their KE and PE as long as the sum remains constant. As they move further apart, PE increases and their speeds must drop so that the KE can decrease.

  • The virial theorem tells how, on average, KE and PE are in balanceBegin with Newtons law of gravity and add an external force FTake the scalar product with x and sum over all stars to getVT is tool for finding masses of star clusters and galaxies where the orbits are not necessarily circular. For system in steady-state (not colliding, etc), use VT to estimate mass

    Assume average motions are isotropic 3r2

    KE (3r2/2) (M/L) Ltot

    Get PE by M = Ltot (M/L) then use galaxy SB to find volume density of stars.

  • In n spatial dimensions, some orbits can be decomposed into n independent periodic motions regular orbits

    Integrals of Motion functions of phase-space coordinates that are constant along any orbit (not time dependent)

    Regular orbits have n isolating integrals and define a surface of 2n-1 dimensionsOrbits in Spherical Potentials at blackboardEquations of motion:2 independent integrals of motion are:

  • VRVR*note that both L and J are used to denote angular momentumEach integral of motion defines a surface in 3-d space (R, VR, V)

    Constant E surface revolves around R-axis

    Constant L surface is a hyperbola in the R, V plane

    Intersection is closed curve and the orbit travels around this curveThe integrals of motion combine (see BT 3.1 for treatment) to produce a differential equation

    where u = 1/R

  • Solutions to this equation have 2 forms:bound = orbits oscillate between finite limits in Runbound = R or u 0

    Each bound orbit is associated with a periodic solution to this equation. Star in this orbit also has a periodic azimuthal motion as it orbits potential center.Relationship between azimuthal and radial periods is:

    is usually not a rational number so orbit is not closed in most spherical potentialsstar never returns to starting point in phase-space

    typical orbit is a rosette and eventually passes every point in annulus between pericenter and apocenter

  • Two special potentials where all bound orbits are closedKeplerian potential point mass

    - radial and azimuthal periods are equal- all stars advance in azimuth by between successive pericenters

    2) Harmonic potential homogeneous sphere

    - radial period is azimuthal period- stars advance in azimuth by between successive pericenters

    Real galaxies are somewhere between the two, so most orbits are rosettes advancing by

    Stars oscillate from apocenter to pericenter and back in a shorter time than is required for one complete azimuthal cycle about center

  • eff = Vo2 ln (R2 + z2/q2) + Lz2/(2R2)(R,z)q= axial ratio

    Resembles of star in oblate spheroid with constant Vc = Voeff rises steeply toward z-axisOrbits in Axisymmetric Potentials at blackboard

  • If only E and Lz constrain motion of star on R,z plane, star should travel everywhere within closed contour of constant eff

    But, stars launched with different initial conditions with same eff follow distinct orbits

    Implies 3rd isolating integral of motion no analytically form

  • Nearly Circular Orbits in Axisymmetric Potentials epicyclic approximationIn disk galaxies, many stars are on nearly circular orbitsderive approximate valid solutions to d2R/dt2 and d2z/dt2.

    Taylor expansion series around (Rg,0) or (x,z) = (0,0)(Ignore higher order terms)Note: x = R Rg yields harmonic potential

    Define two new quantities:Epicyclic and Vertical frequencies

    Then equations of motion become

    x and z evolve like the displacements of 2 harmonic oscillators with frequency and

  • Now relate back to the potentialrecall

    Then the equations become

    Since the angular speed is related to the potential as

    We can now write kappa in terms of the angular speed

    This is related to the (well known) Oort constant BIntegrals of Motion are then

  • The Oort constants, first derived by Jan Oort in 1927, characterize the angular velocity of the Galactic disk near the Sun using observationally determined quantities.A measures shear in the disk would be zero for solid body rotation

    B measures rotation of Galaxy or local L gradientIt can be shown that

    o = A B

    2 = -4B(A-B) = -4Bo at the Sun

    Hipparchos proper motions of nearby stars yield (Feast & Whitlock 1997):

    A = 14.8 0.8 km/s/kpcB = -12.4 0.6 km/s/kpco = 36 10 km/s/kpc

    Sun makes 1.3 oscillations in radial direction in the time to complete one orbit around GCDoes not close (rosette)

  • Continue with Nearly Circular orbit approximation on the board.

    Integrals of motion

    Equations of motion in x , z and y directions

    Derive epicycle shapesX/Y = /(2)

    Pt mass (keplarian rotation curve) = and X/Y=1/2Homogeneous sphere=2 and X/Y=1

  • Orbits in Non-Axisymmetric PotentialsProduce a richer variety of orbits = (x,y) or (x,y,z) cartesian coordinates

    Only 1 classical integral of motion E = v2 + though other integrals of motion may exist for certain potentials which cannot be represented in analytical form

    Orbits in non-axisymmetric potential can be grouped into Orbit Families. Examples can be found in two types of NAPs.Separable PotentialsAll orbits are regular (i.e. the orbits can be decomposed into 2 or 3 independent period motions (in 2 or 3-d)All integrals of motion can be written analyticallyThese are mathematically special and therefore not likely to describe real galaxies. However, numerical simulations for NA galaxy models with central cores have many similarities with separable potentials.

    Distinct families are associated with a set of closed, stable orbits. In 2-d:Oscillates back and forth along major axis (box orbits)Loops around the center (loop orbits)

  • 2-D orbits in non-axisymmetric potentialFor larger R > Rc, orbits are mostly loop orbitsinitial tangential velocity of star determines width of elliptical annulus (similar to way in which width of annulus in AP varies with Lz)Rotation curve is flat with q=1 at large RFor small R
  • box orbit: move along longest (major) axis, parent of familyshort axis tube orbit: loop around minor axis (resemble annular orbit of axisymmetric potentialouter long-axis tube orbit: loop around major axisinner long-axistube orbit: loop around major axisIn 3-d (triaxial potential), there are four families of orbits:Triaxial potentials with cores have orbit families like those in separable potentials.Intermediate and short axis orbits are unstable!Intermediate axis loop orbits are unstable!

  • Scale Free Potentials

    All properties have either a power-law or logarithmic dependence on radius (i.e. ~ r-2)

    These density distributions are similar to central regions of Es and halos of galaxies in general

    If density falls as r-2 or faster, box orbits are replaced by boxletsbox orbits about minor-axis arising from resonance between motion in x and y directions (Miralda-Escude & Schwarzchild 1989)

    Some irregular orbits exist as well (i.e. stochastic motions which wander anywhere permitted by conservation of energy).

  • Stellar Dynamical Systems

    Unlike molecules in a gas, where collisions distribute and average out their motions, stellar systems are governed strictly by gravitation forces.

    For stars, the cumulative effect of small pulls of distant stars is more important than large pulls caused as one star passes close to another. But we will see that even these have little effect over a galaxys lifetime of randomizing or relaxing the stellar motions. Therefore,The smooth Galactic potential of the Milky Way almost entirely dominates the motion of the Sun.

    Consider a system where physical collision are rare. This can be idealized as N point-sized bodies with masses Mi, positions ri and velocities viPotential runs over all pairs twice, hence the 1/2Equations of motion are

  • A general result of the equations of motions is the scalar virial theorem

    where T is KE and U is PE

    since E = T + UTotal mass M and energy E of N-body system define a characteristic velocity and size the virial velocity and virial radius

    The crossing time is a system can be then be definedtc is time scale over which system evolves toward equilibriumtc is constant even for systems far from equilibrium

  • For systems near equilibrium, Vv2 = GM/RvdensityFor systems w/galaxy-like profiles Rv = 2.5 Rh (half-mass radius)

    tc ~ 1.36 (Gh)-0.5 where density is defined within the half-mass radius

    Since crossing time is supposed to be just the typical time scale for orbital motion, we can define the crossing time as

    Under virial assumptions, crossing time depends only on density and increases as density decreases to the square power.

    How does the crossing time relate to the relaxation time, or time it would take the small pulls of distant stars to randomize the stellar orbits?tc = (Gh)-0.5

  • In a distant encounter, the force of one star on another is so weak that stars hardly deviate. We can use the impulse approximation to calculate the forces that a star would feel as it moves along an undisturbed path

    Vt where Vt = 2Gm/(bV)impact parameterBut, when is it a close enough encounter to matter??A strong encounter occurs when, at closest approach, the change in the PE is as great as the initial KE

    Gm2/r 1/2mV2 so r rs = 2Gm/V2 this is the strong encounter radius

    Near the Sun, V ~ 30 km/s, m ~ 0.5 M then, rs ~ 1 AU pretty close!How often does this occur?Assume the Sun is moving with speed V for a time t through a cylinder with radius rs and volume rs2 V t. What is time ts such that n rs2 V t = 1?

    = 1015 yrs with typical solar valuesm

  • Back to considering effects of distant encounters

    Using impulse approximation, a star will have dnenc encounters during a single passage through a systemSurface density of starsarea of annulus with radius b and width dbThe star receives many deflections due to dn encounters, each with random direction, so expected tangential velocity after time t is obtained by adding perturbations in quadratureTotal velocity perturbation acquired in one crossing time

    So a single distant encounter may barely effect the star, but the cumulative effects are important!Each decade between bmin and Rv contribute equally to total deflectionNote: rs=2bmin

  • Now estimate V as the virial velocity Vv where Vv ~ sqrt(GNm/Rv) and let


    .is the total change in a stars velocity per crossing time tcRelaxation time is the time over which the cumulative effects of stellar encounters become comparable to a stars initial velocity

    For galaxies N~1011 stars and tr = 5x108 tc (relaxation important after ~100 million crossings)But, galaxies in general are only ~100 crossing times old cumulative effects of encounters between stars are pretty insignificant!For globular clusters, N~106 or 105 stars and tc~105 yrs tr=5000tc stellar encounters important after ~109-1010 yrsIn denser cores, encounters play a key role

  • Collisionless Dynamics In the continuum limit, stars move in the smooth gravitational field (x,t) of the galaxy. So instead of thinking about motion in 6N dimensions, we can simplify to just 6 dimensions.

    Galaxy may be described by a one-body distribution function (probability density in phase-space):


    -average number of stars in phase-space volume at (x,v) and time t

    Number density (at position x) is then

    n(x,t) = integral f(x,v,t) d3vAverage velocityn(x,t) = integral v f(x,v,t) d3v

    Find equations relating changes in the density and DF as stars move about the galaxy

  • As stars move through a galaxy, how do changes in the density and DF of stars relate to the potential?vxx x+xSimplify to one direction xn(x,t) x is # stars in box between x and x+x at time tAfter time t:

    x[n(x,t + t) n(x,t)] = n(x,t)v(x) t n(x+x,t) v(x+x) tentered in t left in t Left side is the change in # between the two times and right side is change in stars entering and leaving which should be equivalent

  • Take limits at t 0 and x 0

    Equation of continuity stars are not destroyed or addedRate of stars flowing in + rate of stars flowing out is zero. The Collisionless Boltzmann equation is like the EOC, but allows for changes in velocity and relates changes in DF to forces on the stars.Assume acceleration of star dv/dt depends only on potential at (x,t)If dv/dt>0, after t all stars will be moving faster by t(dv/dt)Stars with velocities between v and v-t(dv/dt) move inThose with velocities below v+v have leftThen, the net # of stars that enter the center box after t due to change in v and x

  • In the limit that all s are smallEOC in phase-space spaceUnder gravity, stars acceleration depends only on position

    So, I-D CBE

    And in 3D

    Collisionless Boltzmann Equation the fundamental equation of stellar dynamics

  • Equation holds if stars are neither created or destroyed and change position & velocity smoothly.

    BT describe CBE this way = The flow of stellar phase points through phase space is incompressible; the phase-space density f around the phase point of a given star always remains the same.

    If there are close encounters between stars, these can alter the position and velocity much faster than a smoothed potential. In this case, the effects are given as an extra collisional term on the right side.Since f is a function of seven variables (phase-space and time), the complete solution of CBE is usually too difficult but, velocity moments of CBE can be used to answer specific questions in stellar dynamics

  • Integrate CBE over velocity and apply 0th velocity moment

    1st velocity moment where i=1-3 (3D)

    1. 0th moment of CBEThis is the EOC no surprise since we just integrate over velocity2. 1st moment of CBE

  • From 2nd moment of velocity, velocity dispersion is

    Combine 1 and 2 and divide by n

    3. accelerationkinematic viscositygravitypressure-analogous to Eulers equation in fluid mechanicsEquations 1, 2, and 3 are known as Jeans Equations (Sir James Jeans, 1919) - first applied to stellar dynamics

  • Applying Jeans Equations and CBE Mass Density in the Galactic Disk

    Select tracer stellar type (K dwarfs) and measure density n(z) at height z above disk (coordinates (z, vz) instead of (x, v))Assume potential, DF and number density n do not change with timeAt large z, n(z) 0, thus Eq. 1 gives = 0 everywhereEq. 3 with = z and since =0 we lose the 1st and 2nd termIf we measure how density and sigma changes with z, we get vertical force at any height z.

    Now use Poissons Eq., which relates that force to mass density of the Galaxy.Assume MW is axisymmetric so potential and density only depend on R, zThis is equal to Vc2(R)

  • Since V(R) is ~constant at Sun, let the last term = 0Then, if we know # density wrt z and the velocity dispersion in z, we get density!

    More accurate to determine mass surface density than volume density Oort (1932) measured n(z) for F dwarfs and K giants and obtained (

  • More recent work with fainter K dwarfs (more numerous and evenly spread out) indicates z increases with z:

    z ~ 20 km/s @ 250 pc and 30 km/s @ 1 kpcYields (