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Astron. Astrophys. 358, 759-775 (2000)

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3. High speed stream

In a high speed stream the velocity increases and the density decreases in such a way that the dynamic pressure [FORMULA] remains constant. This means that the solar wind mass flux [FORMULA] decreases, and [FORMULA] increases. When all other parameters remain constant, then it follows from Eq. (5) that neither [FORMULA] nor [FORMULA] should be affected by a high speed stream. But in a HSS the magnetic field is enhanced by a factor of two (see e.g. Gosling et al. 1978, Fig. 3). The enhanced magnetic field transfers more momentum from the ambient solar wind to the central parts of the tail. Thus, the plasma moves faster, it becomes more tenuous. This is the main effect, which dims the comet when it enters an HSS, and which causes the comet to shed off part of its tail.

[FIGURE] Fig. 3. The position of the tail (lower panel) the maximum brightness (middle panel) and the ions per tail length (upper panel) after 0 (solid), 1 (dotted), 2 (dashed), 3 (dash-dotted), 4 (dash-dot-dot-dot) and 5 hours (long dashes).

We started a time dependent calculation from the stationary model 'slow1' (see Table 1). At time t=0 a high speed stream with parameters 'fast1' (see Table 1) first hits the inflow boundary. The solar wind parameters change discontinuously. The discontinuity surface is inclined to the flow direction by 45 o. It is parallel to the interplanetary magnetic field (IMF).

The solar wind discontinuity is not in equilibrium. It resolves according to the Riemann problem into several waves. Since the magnetic field is parallel to the plane of the discontinuity there are only three waves, namely two shocks moving with the fast magnetosonic speed and a contact discontinuity moving with the flow speed. The slow magnetosonic speed and the Alfven speed are both zero since the field is transversal to the direction of wave propagation. Since the magnetic field is small, the resolution of the flow discontinuity can be calculated hydrodynamically. The result, shown in Fig. 1 and in Table 2, can be discribed as follows:


Table 2. Resolution of the discontinuous transition from slow1 to fast1 wind

The component [FORMULA] of the velocity normal to the discontinuity surface increases in two steps. The fast flow pushes from behind onto the slow flow. In the interaction region a density and a pressure hump develop, and the dynamic pressure increases by a factor of 3.3. This interaction region runs through the comet with a velocity of 350 to 470 km/s. It passes the ambient plasma in the computational grid in 30 to 45 minutes.

The central part of the tail reacts much more slowly. Fig. 2 shows the brightness (proportional to the column densities) in the model calculation 0, 1, 2, 3, 4, and 5 hours after the HSS has entered at the lower left corner. The view is perpendicular to the IMF.

The tail is shaped by several effects.

First, the tail turns and so adjusts to the new flow direction of the solar wind. This becomes clearer in the lower panel of Fig. 3 where the position of the tail is plotted at the same instances. (The tail position is defined as the point of maximal brightness in a cross-section perpendicular to the tail.) The tail is first shifted as a whole. But after three hours a kink develops at a distance of 350 000 km behind the nucleus. This kink becomes more pronounced after 4 hours.

Secondly, part of the tail recedes and the new tail adapts to the lower brightness pertinent to the new solar wind conditions. This can be seen most clearly in the middle panel of Fig. 3 where the maximum brightness in sections across the tail is shown. In the first hour, a compressional wave runs through the tail which enhances the density and the maximum brightness all over the tail. But in the next few hours a clear second maximum in the tail develops. A 'cloud' is formed. This cloud moves downward. The tail settles down in a new steady state corresponding to the 'fast1' solar wind condition.

The cloud's motion is accelerated with an acceleration a=5.3 m s-2. For comparison: the velocity along the tail axis can be approximated by an accelerated flow with [FORMULA] for 'slow1', and [FORMULA] m s-2 for 'fast1'. Hence, the acceleration of the cloud is between these two values.

The upper panel of Fig. 3 shows the ions per tail length, i.e., the ions in Fig. 2 integrated over the y-coordinate perpendicular to the tail axis. The curves start for t=0 from a high level. But after 2 hours they consist of a low and a high part, with the transition moving downstream as time proceeds. This shows clearly that the bright tail is disconnected, moves downstream and gives way to a dimmer tail.

A closer analysis shows that part of the tail's mass is severed from the coma and moves downstream. The solar wind has changed its direction. It blows now more from the side. The old tail has a luff and a lee side. This becomes most pronounced at that part of the tail where the severed dense cloud passes by. The massive cloud resists with some success the forces of the wind. On the lee of this cloud the tail is rarified.

The magnetic field is shown in Fig. 4. The field is removed and cannot be replenished by the blocked solar wind (see the 'magnetic hole' in the far tail after 3 and 4 hs in Fig. 4). The cloud itself is compressed by the action of the cross streaming solar wind. At the luff the magnetic field piles up (see the difference between the lower and the upper part of the tail in Fig. 4 after 2, 3 and 4 hs). There is an enhancement of the total pressure in the cloud indicating an adiabatic compression. This confines the cloud and shapes it to the bright bar which connects the two parts of the tail pointing into the new and old flow directions.

[FIGURE] Fig. 4. The magnetic field strength [nT] in the IMF plane. The contour levels are 10(5)50 nT.

Thus, the signature of the action of an HSS is a kink in the tail, formed by a bar connecting two tail segments pointing into different directions. At the inner end of the kink the tail seems to break off.

The kink is observable only when the line of sight forms an angle with the transverse component of the IMF. Fig. 5 shows how the tail would look like for an observer in the plane of the IMF. After three hours a cloud appears in the tail. This cloud marks the position of the kink seen in Fig. 2.

[FIGURE] Fig. 5. The same as Fig. 2 but in a projection onto the perpendicular plane.

As we have argued before, the reduced brightness cannot be attributed to the changes in velocity and density alone, since the dynamic pressure remains unchanged. The effect is to a large extent generated by the enhanced magnetic field, which transfers much more momentum from the bystreaming solar wind into the tail and so accelerates the plasma more efficiently and in this way removes it from the tail which becomes more tenuous. To demonstrate this more convincingly we have repeated the calculation with the same parameters but with unchanged magnetic field. These calculations are shown in Sect. 5.

But the tail morphology also cannot be generated by a simple change of the flow direction. This will be demonstrated in Sect. 4.

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© European Southern Observatory (ESO) 2000

Online publication: June 8, 2000