A study of the wave properties of microscale fluctuations (length scales 01 a. u. and less) in the interplanetary medium using plasma and magnetic field data from Mariner V (Venus 196 7) is presented. The magnetic field data reduction procedure is summarized and descriptions of the MIT plasma data and the combined plasma/field data strips used in the analysis are given. These observations are organized around a model of the solar wind velocity structure.
The large-scale properties of the solar wind are well known (for a comprehensive review, see Hundhausen [1968. Thus, all MHD perturbations are externally convected, regardless of their true direction of propagation in the remaining plasma.
CHAPTER II
Measurements of the field components are returned in digital units (DN) ranging from +511 to -511. Subtraction of the spacecraft field from the XYZ magnetometer measurements yields the true interplanetary magnetic field :J2_. One of the most important steps in data reduction is the choice of time.
A day of magnetic field data plotted using 168.75 sec averages of individual vector measurements. This is the angle of the field outside the solar equatorial plane, and A is a longitudinal angle in this plane measured from the R axis.
CHAPTER III OBSERVATIONS
These authors tentatively concluded that most of the discontinuities in the solar wind are tangential (non-propagating). For the 416 six-hour intervals in the flight with more than 66 percent data return, 33 percent of the. As mentioned above, Alfven waves in the interplanetary medium I. have characteristic association patterns with the large-scale velocity structure of the solar wind.
The best examples of the Alfvn waves (light bars) are found in high-velocity flows and on their trailing edges. Before hour 15, in the most active part of the compression region, the correlations are clear. On days 193 and 194, the amplitude of the field fluctuations has decreased (this period is in the actual high-velocity flow), and the correlation between b and v is extremely good (as in Fig. 2a).
The strength of the magnetic field fluctuations increases on all three axes as we move from the actual high-speed flow into. This correlation is presumably due to conditions in the corona and the radial distribution of the energy. The normal component of the magnetic field has more strength than the other components during the latter half of day 236.
51 represents the amount of power in the very high frequency fluctuations of the magnetic field in each three-hour interval;. The purest examples of outward-propagating Alfvén oscillations are found in high-velocity flows and at their trailing edges. Preliminary power spectra and cross spectra of the interplanetary magnetic field in the frequency range from 1/(107.5 min.) to.
In the previous section dealing with wave spectra, we made no reference to or use of the vector nature of magnetic field fluctuations. We first discuss the general properties of microscale variations in the magnetic field and then examine in detail specific periods of dynamical interest.
CHAPTER IV
There is also no convincing reason to expect that such collisions will generate Alfve'n waves that propagate outward exclusively in the rest frame of the wind. It has been suggested [Parker, 1969] that the energy equation of the solar wind near the Sun is dominated by the transport and deposition of energy by waves. It seems very likely that the exclusively outwardly propagating Alfve'n waves found in the main bodies of solar wind streams are generated in the vicinity of the Sun, far away from sharp changes in velocity.
The appearance of larger amplitude, purely outwardly propagating Alfve waves in the main bodies of the high velocity. Rather, it is mainly due to the fact that the energy input to the base of the solar wind for the slow gas that precedes the fast is such that it produces higher densities in the slow gas (as well as lower velocities and temperatures) at l a. If they are polarized in the solar equatorial plane, they will be partially converted into magnetoacoustic waves due to the sudden turn.
It appears that all the observed properties of waves in the solar wind follow naturally from this model. Most of the waves in the interplanetary medium are Alfve'n waves propagating outward from near the Sun, and are remnants of the processes that provide energy to drive the various solar wind streams. The main points to be explained are that the minimum component of the fluctuations is parallel to !:.B• the.
96- . direction of the mean magnetic field, and that the maximum component is in the direction ~Bx~, where ~ can be regarded as the direction of the solar wind flow. There appears to be no plausible anisotropy in the generation of the Alfve'n waves that would favor the .::_Bx~ direction. As the waves are convected outward, the sun's rotation causes the mean field to rotate in the plane perpendicular to the moment the spiral is generated.
CHAPTER V
Inserting equations (17) and (24) into equation (8) and averaging over a wavelength of the waves gives. The second is the kinetic energy density of waves convected by the bulk velocity, and the third is the radial component of the Poynting vector. For this value of U, the solution for U(Z) given by equation .. Z ) passes through the critical point, and this is the sun.
It is readily seen that the inclusion of the wave energy flux in Parker's fluid polytropic model of the solar wind [Parker, 1963 J. The flow velocity of the solar wind at this level is generally much smaller than the Alfve'n velocity, so that the outward energy flux along the reference plane in Alfve'n coronal waves is given by. These flux estimates should be compared with energy flux estimates due to thermal conductivity from the lower corona on the order of 2 x 1027 ergs/.
So the remaining free parameters for these calculations are E, which is representative of the wave energy flux at r , and T , the coronal temperature there. The parameter H as defined above (Equation 31) ) is the negative of the gravitational potential energy per gas atom in. The effect of the waves (increasing E from zero) is to increase the speed of the solar wind at the reference point r and decrease the critical one.
This is also due to the specific form of interaction between wind and waves in this model. The second term in equation (56) is the increase in the square of the solar wind speed at infinity due to the waves. This property of wave/wind interaction qualitatively explains why the additional velocity at infinity due to waves depends directly on the VA0/V ratio.
For the E=10 case, the density profile is almost the same as the static profile, except at the top of the static atmosphere. This is not unexpected, since in this model the purely static solutions have Alfvdn velocities that approach infinity as the top of the atmosphere is reached. With the wave-driven solutions we see that the opposite can be true; that is, higher velocities with lower densities at 1 a.
Regarding the inclusion of the wave energy fluxes, there are two questionable assumptions in our representation. As long as oB/B < 1 this is not unreasonable, since both from a theoretical and observational point of view, Alfv waves in the solar wind do not appear to damp rapidly under this condition. In our model, the wave amplitudes of the coronal Alfve'n waves generally increase as r increases, so that oB/B can equal 1 at some point.
If this point occurs at distances much larger than the critical radius rC'', the large change in the solutions due to wave damping will occur. If the damping has a large effect at less than the critical radius, the bulk properties of the solutions far from the Sun may change significantly, and a more complete theory is needed. In all numerical solutions discussed above, oB/B only becomes comparable to 1 when values of r are greater than.
The correction of this defect in the model will likely require a detailed theory of the actual mechanisms generating the waves in the lower corona, as well as a complex integration scheme. With these shortcomings in mind, we believe that the model is a reasonable first attempt to determine the effects of coronal waves on wind dynamics. The second is the fact that wave-driven solutions can easily produce the high speeds and low densities often observed in the solar wind at 1 a.
CHAPTER VI
The detailed observational behavior of the waves far outside the orbit of the earth (where the magnetic field is tightly wound in the classical spiral) is of particular interest. Although the model is unrealistic in that it assumes a single fluid plasma with no wave damping, it provides useful insight into the qualitative effects of the observed waves on the actual solar wind. The next step is a two-fluid model with a more sophisticated treatment of the energy transport equations (including wave damping).
The effects of the rotation of the sun and the spiral field patterns on wave-driven solutions are also of great importance. A more fundamental problem is the nature of the processes that generate the waves at or near the sun. Finally, the thesis resolved many observational points and gave some consideration to the theoretical points these observations raise.
