8. TIME DEPENDENCE 138 where the integral is taken over the entire star. In a convective atmosphere, P ∝ργ, so ρ ∝ P1/γ. If we make the thin atmosphere approximation throughout the star, just to gain an order-of-magnitude estimate, thenP =gΣ, and so
τadj ≈g−1F−1/3
Z
ρ−2/3dP =g−1F−1/3ρ−2/30
Z P P0
−2/3γ
dP (8.27)
= 1
1− 3γ2 g−1F−1/3ρ−2/30 P02/3γPf1−2/3γ−P01−2/3γ. (8.28) Now γ is typically 5/3 outside of the ionization zone, so 2/3γ=2/5, and hence
τadj ≈ 5P02/5 3gF1/3ρ2/30
Pf3/5−P03/5≈ 5P02/5Pf3/5
3gF1/3ρ2/30 = 5v2s,0 3γgvc,0
Pf P0
3/5
= v2s,0 gvc,0
Pf P0
3/5
, (8.29) where we have made the approximation that the core pressure vastly exceeds the pressure at the top of the convection zone. Nowvc,0 is typically aroundvs,0/10, and vs,0 is typically around 106cm/s≈100sg, so the prefactor is around 103s. Typically P0 ∼ 105erg/cm3, and Pf ∼ gM/(4πR2) ∼ 1015erg/cm3 for the sun, so the overall timescale is around 103+6 = 109s, scaling roughly as M6/5R−12/5. For the fully convective star considered in simulation, M = 0.3M and R = 2.65R, so this timescale is smaller by a factor of 25, giving around 4×107s, or roughly a year.
On the other hand, the core adjusts its temperature in time τcore = mcorecpTcore
Le =fMLin
Le τK =fMGM2
2RLe, (8.30)
where fM ≈1/10 is the fraction of the star’s mass in the core andτK is the Kelvin timescale for the star. This is typically of order ten million years, and so ifLe =Lin the core’s adjustment timescale is of order a million years. As this is much shorter than the timescale required to maintain adiabaticity, the star may be approximated as being adiabatic at all times after the cooling wavefront reaches the core.
8. TIME DEPENDENCE 139 in these stars, meaning that the bloating effect is further decreased by the amount given in figure 2.7.
To investigate these effects in the transient case, we first simulated a sun-type star initially illuminated by Le = L and watched as the illumination was turned off. The results of this are shown in figure 8.7. The key feature we see here is that the luminosity sits fixed near the initial steady-state value at the surface, and that the main effect of time evolution is to push the transition between this value and the nuclear-burning value deeper into the star. The depth at which this occurs is between Σ = 104g/cm2 and Σ = 105g/cm2, just slightly deeper than the point in our steady-state calculations where the radiative-convective transition arises in illuminated equilibrium in this sort of star. This feature is not unique to stars of M =M. Figure 8.8 shows a star of the form examined in the preceding section, but withLe = 10L. This star exhibits a similar transition between radiative and convective heat transport in the steady state and hence exhibits a similar transient adjustment process. The story behind the evolution of stars such as these is then that the external illumination shuts off convection beyond a certain depth. When the illumination is removed, that radiative region dampens the resulting change in temperature exponentially into the star, while the convective region maintains a luminosity close to the initial steady-state value. This is precisely what we see, but we can further test this notion by examining the star on longer timescales. If this story is correct, the star will slowly turn the radiative zone back into a convection zone, and in the process the luminosity profile will settle down to have L=L everywhere.
To determine if this is the case, the simulation was run for another 108s and found indeed to be out of equilibrium, a feature not seen in any of the previous scenarios considered. The results of the longer simulation are shown in figure 8.9. Note that as the luminosity transition region pushes deeper into the star, the magnitude of the transition falls. Over even longer timescales, the equilibration continues but slows down somewhat, as shown in figure 8.10. The adjustment time for this process is on the order of the thermal timescale for the entire region in which the mode of heat transport shifted from being convective to being radiative, perhaps decreased by a factor of 10 to account for the relatively small temperature changes required to do this at high Σ. As a result, the full adjustment process requires timescales beyond the realm of validity of our lower boundary condition on T. Fortunately all that matters for our purposes are the trend and timescale involved, which are clearly seen in the simulations which are accessible.
8. TIME DEPENDENCE 140
0 1 2 3 4 5 6 7 8
−0.35
−0.30
−0.25
−0.20
−0.15
−0.10
−0.05 0.00 0.05
∆T/T
Lin=Lsun,∆t =1e8s
0 1 2 3 4 5 6 7 8
Log Σ 1.0
1.2 1.4 1.6 1.8 2.0 2.2
L/Lsun
t=0e7s t=1e7s t=2e7s t=3e7s t=4e7s t=5e7s t=6e7s t=7e7s t=8e7s t=9e7s t=10e7s
Figure 8.7: ∆T /T0 (top) andL/L(bottom) versus log Σ (in g/cm2) for a star of mass M, radius R, and luminosity L. The external heat was put in at Σ = 103g/cm2 and linearly decreased fromL to zero over the course of 108s. Color represents time, with the simulation beginning at violet and ending with red.
8. TIME DEPENDENCE 141
0 1 2 3 4 5 6 7
−0.7
−0.6
−0.5
−0.4
−0.3
−0.2
−0.1 0.0 0.1
∆T/T
Lin=0.1Lsun,∆t =1e8s
0 1 2 3 4 5 6 7
Log Σ 0
2 4 6 8 10 12
L/Lsun
t=0e7s t=1e7s t=2e7s t=3e7s t=4e7s t=5e7s t=6e7s t=7e7s t=8e7s t=9e7s t=10e7s
Figure 8.8: ∆T /T0 (top) and L/L (bottom) versus log Σ (in g/cm2) for a star of mass 0.3M, radius 2.65R, and luminosity 0.1L. The external heat was put in at Σ = 103g/cm2 and linearly decreased from 10L to zero over the course of 108s.
The simulation was then run for an additional 108s with no external heating. Color represents time, with the simulation beginning at violet and ending with red.
8. TIME DEPENDENCE 142
0 1 2 3 4 5 6 7 8
−0.5
−0.4
−0.3
−0.2
−0.1 0.0 0.1
∆T/T
Lin=Lsun,∆t =2e8s
0 1 2 3 4 5 6 7 8
Log Σ 1.0
1.2 1.4 1.6 1.8 2.0 2.2
L/Lsun
t=0e7s t=2e7s t=4e7s t=6e7s t=8e7s t=10e7s t=12e7s t=14e7s t=16e7s t=18e7s t=20e7s
Figure 8.9: ∆T /T0 (top) andL/L(bottom) versus log Σ (in g/cm2) for a star of mass M, radius R, and luminosity L. The external heat was put in at Σ = 103g/cm2 and linearly decreased from L to zero over the course of 108s. It was then run for another 108s at that value. Color represents time, with the simulation beginning at violet and ending with red.
8. TIME DEPENDENCE 143
0 1 2 3 4 5 6 7 8
−0.6
−0.5
−0.4
−0.3
−0.2
−0.1 0.0 0.1
∆T/T
Lin=Lsun,∆t =1e9s
0 1 2 3 4 5 6 7 8
Log Σ 1.0
1.2 1.4 1.6 1.8 2.0 2.2
L/Lsun
t=0e8s t=1e8s t=2e8s t=3e8s t=4e8s t=5e8s t=6e8s t=7e8s t=8e8s t=9e8s t=10e8s
Figure 8.10: ∆T /T0 (top) and L/L (bottom) versus log Σ (in g/cm2) for a star of mass M, radius R, and luminosity L. The external heat was put in at Σ = 103g/cm2 and immediately decreased from L to zero over the course of 108s before being run for another 109s. Color represents time, with the simulation beginning at violet and ending with red.
Part II
Applications in Astronomy
144
9
X-Ray Binaries
You know, you blow up one sun and suddenly everyone expects you to walk on water.
– Lt Col. Samantha Carter, Stargate SG-1 Season 8 Episode 17 The results presented thus far have been general, in the sense that while there were motivating examples of phenomena of interest, many avenues were pursued to provide a picture of the phenomenology of pulsar-companion systems. We are now interested in examining the specific case in which the pulsar interacts with its companion to produce transient X-ray emissions. This case has long been studied1, though conclusions have proven scarce. In addition, while in previous chapters the companion was a passive agent, here we will consider the role it plays in influencing its own fate. The first section deals with the isotropic illumination case, while the second discusses the effects of anisotropy.
9.1 Accretion rate
The initial heating of the star causes it to expand at some rate ˙R. This rate is everywhere the same in the atmosphere due to the expansion being driven by deep heating, as discussed earlier. As the atmosphere of the star falls off exponentially
1J. C. Brown and C. B. Boyle. “An exploratory eccentric orbit ’Roche lobe’ overflow model for recurrent X-ray transients”. In: Astronomy and Astrophysics 141 (Dec. 1984), pp. 369–375;
H. Ritter, Z.-Y. Zhang, and U. Kolb. “Irradiation and mass transfer in low-mass compact binaries”.
In: Astronomy and Astrophysics 360 (Aug. 2000), p. 969. eprint: astro-ph/0005480; A. R. King et al. “Mass Transfer Cycles in Close Binaries with Evolved Companions”. In: The Astrophysical Journal482 (June 1997), pp. 919–928. eprint: astro-ph/9701206.
145
9. X-RAY BINARIES 146 in the radial coordinate above the photosphere, no significant accretion is expected until this region approaches the Roche lobe radiusRb. The accretion rate is expected to be2
M˙ =√
2πRhsvsρ(Rb). (9.1)
Herehs ≈Rvs2/v02, wherev0 is the orbital velocity of the star. This is due to the fact that in the vicinity of the Roche lobe, the pressure profile is set by orbital parameters rather than the thermal structure of the star. As a result, we may write
M˙ ≈√
2πR2v−20 v3sρ(Rb). (9.2) If the accretion rate is low3, it typically means that ρ is low at Rb, the Roche radius, and hence that we are in the upper portion of the atmosphere. This allows us to make use ofρ∝exp(−r/hs) and write
M˙ ≈√
2πR2v−20 vs3ρ0exp rv20 Rvs2
!
=√
2πΩ−2v3sρ0exp rR20Ω2 Rv2s
!
, (9.3) where ρ0 is chosen to make this relation true and r is a Lagrangian quantity. For the accretion to be significant we must have R ≈ Rb, for R, Rb hs because v0 ≈107cm/s105cm/s≈vs. Thus
M˙ ≈√
2πΩ−2vs3ρ0exp rR02Ω2 Rbvs2
!
. (9.4)
In Part 1, we found that only stars with deep convection can swell to the point where R∼Rb, so we restrict ourselves to stars of this form. As a result,M < 1.2M. Using Mp ≈2M, we may approximate4
Rb ≈0.46R0 M M +Mp
!1/3
, (9.5)
yielding
M˙ ≈√
2πΩ−2vs3ρ0exp 2rR0Ω2(M+Mp)2/3 M2/3v2s
!
, (9.6)
2Brown and Boyle, op. cit.
3Using Eq. (9.2), we find that ˙M ∼1024ρcm3/s. Based on the data in Appendix E, the exponential atmosphere assumption holds at least up toρ∼10−8g/cm3, so we are safe making this assumption if ˙M <1016g/s. As will become clear subsequently, this is much larger than the typical values we will encounter.
4B. Paczyński. “Evolutionary Processes in Close Binary Systems”. In: Annual Review of Astronomy and Astrophysics9 (1971), p. 183. doi: 10.1146/annurev.aa.09.090171.001151.
9. X-RAY BINARIES 147 where
R0 = G(M +Mp) Ω2
!1/3
. (9.7)
Typical atmospheric temperatures are such that µ=mp, so γ = 5/3 and v2s = 5kBT
3mp . (9.8)
Thus we can compute all of the quantities in the exponential.
Now we haven’t yet fixed r or ρ0, and so we actually have the freedom to absorb any constants we wish. Furthermore, relative to the exponential the dependence on T is negligible, so we may let T →T0 for some reference photospheric temperature T0 and absorb it as well. Thus we will write instead
M˙ ≈exp 2rR0Ω2(M +Mp)2/3 M2/3vs2
!
. (9.9)
We now no longer have the freedom to pick the zero-point ofr. Rather, it is uniquely determined given ˙M at some time. Without solving for it, though, we may write
M¨ = 2R0Ω2(M +Mp)2/3
M2/3vs2 r˙M .˙ (9.10)
Using ˙r= ˙R and dividing through by ˙M yields
∂tln ˙M = 2rR0Ω2(M +Mp)2/3 M2/3vs2
R.˙ (9.11)
This equation is independent of the zero-point of r, forr no longer appears anywhere in it. Given ln ˙M at some point in time, we may use this relation to determine it at any subsequent point so long as we know ˙R.