The strahl represents one of the four populations that make up the electron solar wind. Unlike the thermal, isotropic core, the suprathermal halo and the high-energy super-halo , all of which follow a radial trajectory away from the sun, the strahl flows parallel or antiparallel to the ambient solar wind magnetic field, i.e., it is field-aligned. The general consensus is that the solar wind originates in the solar corona through a mixture of Coulomb scattering and interactions with broadband turbulence ; however, there is still no clear consensus on the specific mechanisms involved in the formation of any of its four individual populations.

The strahl is thought to be formed from the initial high-energy electron tail in the absence of Coulomb collisions and is focused along the strong magnetic field although recently have put forward the possibility that the strahl is formed directly from the halo by pitch angle scattering off of whistler waves. In the absence of scattering, the mirror force will focus the strahl in pitch angle as it propagates away from the sun. This strong focusing is, however, not observed. Observations show that the strahl actually broadens with radial distance , beginning about where scattering would be expected to dominate over focusing (0.5 AU) . Several suggestions as to the free-energy sources available to drive the broadening have been put forth (e.g., see ). The possible mechanisms include scattering of strahl electrons through interactions with sunward propagating whistler waves (e.g, see ), broadband whistler turbulence and Langmuir waves .

and have shown that the strahl and halo densities vary inversely with radial distance from the sun, implying that at least a portion of the strahl may be degraded in energy and integrated into the halo. The processes active to drive this are not known, but it has been suggested that the same processes responsible for broadening the strahl may also be capable of scattering strahl electrons into the halo. This is envisioned as being a slow and continual erosion of the strahl and a corresponding buildup of the halo through inelastic scattering.

The characteristics of the strahl in slow wind have been studied by , , and . They find similar strahl profiles in the fast and slow winds with perhaps a slightly higher density in the fast wind and increased broadening during periods of slow wind. in a large statistical study claimed that there was no strahl in about 25 % of the data. This was an automated statistically based study, and analysis is performed only during times when specific criteria are met. They mention no long contiguous time periods of no strahl or any correlation of the absence of the strahl with solar wind speed.

have shown observations of what appears to be strong local diffusion of the strahl within the energy range where the strahl and halo overlap, suggesting that the buildup of the halo may occur rapidly within regions where the intense disruption of the lower edge of the strahl occurs. Again, it is not certain what processes are responsible for the observations.

In this paper we will detail a set of initial observations made at 1 AU by the Plasma Electron And Current Experiment (PEACE) instrument on the Cluster spacecraft. The observations show that during time periods of slow wind (≲ 425 km s-1), the strahl is often not seen. In place of the strahl, there is an unstructured electron plasma that is smeared in azimuth and elevation angle. This unstructured distribution may be a consequence of counting statistics due to low count rates and is possibly coupled with intense scattering.

This study makes use of data from a number of experiments on board the Cluster spacecraft, including PEACE , the Fluxgate Magnetometer (FGM) , the Electric Field and Waves (EFW) experiment , the Waves of High frequency and Sounder for Probing of Electron density by Relaxation (WHISPER) , and the Spatio-Temporal Analysis of Field Fluctuations (STAFF) experiment .

The primary data used in the study come from PEACE. PEACE consists of two hemispherical electrostatic analyzers designated HEEA (high-energy electrostatic analyzer) and LEEA (low-energy electrostatic analyzer). These are separated by 180∘ on the satellite and differ only in their geometric factors (HEEA's geometric factor is larger than LEEA's). Despite their names, both can cover the energy range 0.6 eV to 26 keV. The analyzers' fields of view are perpendicular to the spacecraft spin axis, which is about 5∘ off GSE -Z, and cover 180∘ in elevation in 12 sectors. The full 360∘ of azimuth is covered in one rotation of the spacecraft so that a three-dimensional snapshot of the electron distribution is accumulated once per spin (∼ 4 s).

Because of telemetry restrictions, PEACE generally returns only a subset of the total data collected onboard. Exactly what is returned depends on the instrument mode, which can be separately commanded for each analyzer on each of the four spacecraft. The telemetry rate, as well as the amount of data being returned, determines the time cadence at which full three-dimensional distributions are downloaded. During the time intervals used in this paper, all satellites were operating in burst-mode telemetry and PEACE was returning one 3-D distribution per spin. The data were being returned in 6 or 12 elevations bins with 30 or 60 energy steps and 16 or 32 azimuths, depending on the data allocation available.

PEACE data are used to characterize the electron plasma through both moments and visualization tools that allow various aspects of the electron velocity distribution function (eVDF) morphology in velocity space to be highlighted. The FGM five vector per second data are used to characterize the local magnetic field, while the full-resolution FGM data together with the STAFF waveform data are used to construct magnetic field power spectra. Both the EFW and WHISPER are used in the formation of the electron moments; EFW provides the spacecraft potential used to correct the measured electron energy, and WHISPER provides the flags necessary to filter out times during which the computed moments may be contaminated by perturbations created by WHISPER active sounding.

ϕ–θ plots are used throughout this paper to illustrate features in the eVDFs. This mode of presentation allows one to show the entire three-dimensional distribution function at a given energy as a two-dimensional projection. The three columns of ϕ–θ plots in Fig.  serve as examples of some of the features available in the plot format. Each set of ϕ–θ plots uses data from the same eVDF and shows a contiguous set of PEACE energy shells from 30.1 to 110.3 eV. To fully characterize an eVDF requires one ϕ–θ plot per returned energy step; however, in general any feature of interest can be characterized by a small subset of the returned energy steps. The data are plotted as a function of ϕ (the spacecraft spin angle) and θ (the instrument elevation angle). For Cluster, the instrument frame of reference is offset by about 5∘ from the -Z GSE frame of reference so that the ecliptic plane makes an angle of about -95∘ with respect to the spacecraft spin axis (0∘ being parallel to the spin axis). A rotation angle of 0∘ is the position at which the instrument aperture lies in the plane containing the sun and the spacecraft x axis. Plots are logarithmically scaled in intensity. Individual plots may be auto-scaled as is the left-hand column of plots in Fig.  or assigned a common scaling range as in the center and right-hand column of plots. When auto-scaling is used, the range is shown above each plot, while when using a common scaling range the scaling is given in a color bar at the bottom of the column. The energy associated with each individual plot is also shown above it in the following format: raw center energy–potential corrected center energy. The dot and triangle in the plots are the projections of the average magnetic field head and tail, respectively.

The three columns of ϕ–θ plots in Fig.  make use of different display options. The eVDF displayed is typical of that seen in the solar wind and the strahl and core/halo populations are explicitly labeled in the left set of plots. The core and halo both flow radially away from the sun and overlap significantly in phase space, which makes it impossible to separate the two populations in a ϕ–θ plot. For this reason we refer to the two populations within a ϕ–θ plot by the single name core/halo. The strahl is field-aligned and approximately centered on one of the magnetic field projection points depending on whether the flow is parallel or antiparallel to B. When the magnetic field has an large off-radial component, the strahl will be shifted off the core/halo as it is in Fig. . Even with a small off-radial component, the shift in the strahl will be evident although there still may be a considerable overlap of the populations. The first and last two columns of plots in the figure illustrate different scaling formats. The left-hand column of plots is auto-scaled, while the center and right-hand columns use a common scaling for all plots. The right-hand column of plots shows the data plotted in units of counts per accumulation (cnts acc-1) with all smoothing and contouring turned off. Here each grid cell in the plot exactly matches the instrument angular resolution.

The basic solar wind or foreshock eVDF can tentatively be broken in energy such that below the break point the core/halo form the dominant electron population and above the break point the strahl (and possibly any return foreshock electrons) are the dominant population(s). The energy range above the break point will be defined as the nominal strahl energy range in this paper. In Fig.  the break point occurs at about 70.5 eV. It is common to find that when the solar wind speed is less than 425 km s-1, there is no strahl population observed above the break point and that any electrons located above this energy appear to be disorganized and diffusely spread in angle. The disorganization is most directly seen in the ϕ–θ plots but can also be inferred through plots of the GSE spherical electron fluid velocity angles determined at each returned energy step.

We have identified 37 intervals from the 181 we analyzed that have no observable strahl. In this paper we will look in depth at two such periods, presenting the basic characteristics observed in the plasma and contrasting them to what is observed in typical solar wind. The features illustrated in the two events are common to one degree or another in all 37 events where no strahl was observed. The first interval spans the time from 09:55 to 12:55 UT on 17 January 2009 and the second that from 22:15 to 23:43 on 28 January 2011. During both events PEACE was returning data in 60 energy × 16 azimuth × 6 elevation bins. We used data from only a single spacecraft in analyzing each event. Generally this was C2 unless it was not returning data, in which case data from C3 or C4 were used. In 2011 C2 lost one of its extreme elevation sensors, and thereafter the data used came from either C3 or C4.

We have looked at every interval between 2001 and 2012, inclusive, when the spacecraft are both upstream of the bow shock and returning data using burst-mode telemetry. There are 181 such intervals, of which 37 contained times when the strahl was absent. The absence of the strahl appears to have some dependence on the solar wind speed, occurring primarily during periods of slow wind (≲ 425 km s-1). The probability of occurrence appears to increase with decreasing wind speed.

Figure illustrates the difference in the observed solar wind speed and temperature anisotropy during periods when the strahl is present and absent. The top two plots in the figure show probability distribution functions (PDFs) of the solar wind speed and temperature anisotropy during the 37 periods where the strahl is absent, while the bottom two plots show the same for the 144 periods when the strahl is present. The PDFs are normalized to area (number of events in each bin divided by the total number of events in the plot) so that the sum of the PDF values in any plot is 1.0. Averages of the quantities in each plot are shown either next to the right or left y axis. PDFs of the speed and anisotropy data when the strahl is present are basically what is observed in the majority of pure solar wind events (see, e.g., ). The spread in the temperature anisotropy is a consequence of the presence of the field-aligned strahl, which is often observed to have temperature ratios greater than 1 and as high as 3 . Intervals when there is no strahl occur almost exclusively at lower solar wind speeds. Such intervals have a temperature anisotropy close to 1 (the approximate temperature anisotropy associated with the core/halo).

Why the strahl is absent is not clear. The primary reason for the uncertainty is that the current set of observations are not taken in regions where the actual disruption of the strahl is taking place. Without knowing where this occurs and under what plasma conditions it occurs, it is almost impossible to begin to speculate on the mechanisms behind it. Given what little information we have, however, two of the more likely explanations are scattering through intense broadband whistler turbulence and the possibility that there was no initial formation of the strahl in the inception of the solar wind.

We have looked at the magnetic field spectral power during times when there is no observable strahl, an example of which is shown in Fig. . The spectrum is a combination of spectra built from the full-resolution FGM data (blue portion of the total spectrum) and the STAFF waveform data (orange portion of the spectrum), which are normalized at their overlap between 1 and 2 Hz. There is no indication of any local intensification(s) in the power in the STAFF portion of the spectrum, which might indicate the presence of intense local turbulence. If we are using this as a signature of scattering, then this is not locally active or has already taken place farther upstream. It should be noted that foreshock spectra in the presence of a strahl signature do tend to show broadband intensity enhancements in the STAFF portion of the spectrum (between 10 and 100 Hz in this figure). These enhancements are more than likely evidence of whistler or other broadband turbulence driven by the free energy in the counterstreaming return and strahl electron populations in the foreshock.

The problem with using scattering to fully disrupt the strahl is that it is generally not intense enough and provides no obviously viable means to explain the observations of side-by-side regions of unstructured and structured solar wind eVDFs separated by narrow transition regions. These observations are puzzling, suggesting that other factors are active in addition to solar wind speed. Figures  and show that the transition from a plasma with a strahl to one without occurs very rapidly and often without any significant difference in wind speed. A detailed example of one transition is given in Fig.  This shows consecutive plots of the 88.3 eV energy step across the last transition from a strahl-dominated to a non-strahl plasma in Fig. . Between 23:36:52 and 23:36:56 UT, the strahl drops out of the eVDF with no effect seen in the return population. The existence of the return population across the time period means that there is no transition between the foreshock and solar wind. It is difficult to understand how these boundaries are set up and maintained.

A second possibility is that the strahl is not uniformly produced in the corona with the solar wind. In this scenario the mixture of regions at 1 AU with a presence or absence of the strahl might be the result of the mixing of flux tubes (viz., ). This would provide for the narrow transition boundaries seen between the two plasma regimes. The supposition would place the flux tube boundary associated with the change in plasma state in Fig.  at about 23:56 UT. Figure  shows a plot of the electron plasma data and magnetic field. The plot is formatted almost identically to the plot shown in Fig.  but covers only a narrow slice of time centered on the flux tube crossing. In this figure, the data are gridded at 0.2 s, which is the native time resolution of the magnetic field data used, while Fig.  is gridded at 4.0 s, which is approximately the time resolution of the plasma data. Crossing the flux tube, there is a noticeable shift in both the average electron temperature and fluid velocity. Both of these can be traced to the existence or absence of the strahl. There are only very minor changes in the magnetic field across the flux tubes. This is no azimuthal rotation and only very small changes in the average elevation angle and average magnitude. Note that the apparent large azimuth rotation in Fig.  at this time is a rotation in flow as the strahl drops out of the eVDF and does not represent a rotation in the magnetic field.

In periods associated with slow wind (≲ 425 km s-1), there is often, but not always, no strahl observed in the solar wind or in the foreshock. It is not clear why the strahl is not seen in the eVDF during those slow solar wind intervals. Is the strahl simply not formed within the conditions under which the slow wind is set up or is it somehow disrupted at some time during or after its formation? It should be emphasized that these suppositions are based on what maybe considered a weak statistical study. There are insufficient observations during times when Cluster is upstream of the shock and returning data in burst-mode to accumulate a sufficient number of events to allow a full statistical analysis of the results. A statistical analysis will be presented in a follow-up paper if we are able to show that a similar analysis can be sufficiently carried out using the lower-resolution data that are commonly used in the solar wind.

With the exception of PEACE data, which were obtained from the Mullard Space Science Laboratory (MSSL) science data archive (http://www.mssl.ucl.ac.uk/missions/cluster/about_peace_data.php), all data were obtained from the Cluster Science Archive (CSA, http://www.cosmos.esa.int/web/csa).