scholarly journals A physical model for the magnetosphere of Uranus at solstice time

2020 ◽  
Vol 643 ◽  
pp. A144
Author(s):  
Filippo Pantellini
Keyword(s):  
Magnetic Field ◽  
Solar Wind ◽  
Magnetic Dipole ◽  
Rotation Axis ◽  
Double Helix ◽  
Mass Density ◽  
Fluid Velocity ◽  
Rotation Period ◽  
Magnetospheric Tail ◽  
Field Lines

Context. Uranus is the only planet in the Solar System whose rotation axis and orbital plane are nearly parallel to each other. Uranus is also the planet with the largest angle between the rotation axis and the direction of its magnetic dipole (roughly 59°). Consequently, the shape and structure of its magnetospheric tail is very different to those of all other planets in whichever season one may consider. The only in situ measurements were obtained in January 1986 during a flyby of the Voyager II spacecraft. At that date, Uranus was near solstice time, but unfortunately the data collected by the spacecraft were much too sparse to allow for a clear view of the structure and dynamics of its extended magnetospheric tail. Later numerical simulations revealed that the magnetic tail of Uranus at solstice time is helically shaped with a characteristic pitch of the order of 1000 planetary radii. Aims. We aim to propose a magnetohydrodynamic model for the magnetic tail of Uranus at solstice time. Methods. We constructed our model based on a symmetrised version of the Uranian system by assuming an exact alignment of the solar wind and the planetary rotation axis and an angle of 90° between the planetary magnetic dipole and the rotation axis. We do also postulate that the impinging solar wind is steady and unmagnetised, which implies that the magnetosphere is quasi-steady in the rotating planetary frame and that there is no magnetic reconnection at the magnetopause. Results. One of the main conclusions is that all magnetic field lines forming the extended magnetic tail follow the same qualitative evolution from the time of their emergence through the planet’s surface and the time of their late evolution after having been stretched and twisted several times downstream of the planet. In the planetary frame, these field lines move on magnetic surfaces that wind up to form a tornado-shaped vortex with two foot points on the planet (one in each magnetic hemisphere). The centre of the vortex (the eye of the tornado) is a simple double helix with a helical pitch (along the symmetry axis z) λ = τ[vz+Bz/(μ0ρ)1/2], where τ is the rotation period of the planet, μ0 the permeability of vacuum, ρ the mass density, vz the fluid velocity, and Bz the magnetic field where all quantities have to be evaluated locally at the centre of the vortex. In summary, in the planetary frame, the motion of a typical magnetic field of the extended Uranian magnetic tail is a vortical motion, which asymptotically converges towards the single double helix, regardless of the line’s emergence point on the planetary surface.

scholarly journals Magnetohydrodynamic simulations of a Uranus-at-equinox type rotating magnetosphere

2020 ◽  
Vol 633 ◽  
pp. A87 ◽  
Author(s):  
L. Griton ◽  
F. Pantellini
Keyword(s):  
Magnetic Field ◽  
Solar Wind ◽  
Supersonic Flow ◽  
Magnetic Dipole ◽  
Stellar Wind ◽  
Rotation Axis ◽  
Spin Axis ◽  
Mhd Equations ◽  
Field Lines ◽  
The Magnetic Field

Context. As proven by measurements at Uranus and Neptune, the magnetic dipole axis and planetary spin axis can be off by a large angle exceeding 45°. The magnetosphere of such an (exo-)planet is highly variable over a one-day period and it does potentially exhibit a complex magnetic tail structure. The dynamics and shape of rotating magnetospheres do obviously depend on the planet’s characteristics but also, and very substantially, on the orientation of the planetary spin axis with respect to the impinging, generally highly supersonic, stellar wind. Aims. On its orbit around the Sun, the orientation of Uranus’ spin axis with respect to the solar wind changes from quasi-perpendicular (solstice) to quasi-parallel (equinox). In this paper, we simulate the magnetosphere of a fictitious Uranus-like planet plunged in a supersonic plasma (the stellar wind) at equinox. A simulation with zero wind velocity is also presented in order to help disentangle the effects of the rotation from the effects of the supersonic wind in the structuring of the planetary magnetic tail. Methods. The ideal magnetohydrodynamic (MHD) equations in conservative form are integrated on a structured spherical grid using the Message-Passing Interface-Adaptive Mesh Refinement Versatile Advection Code (MPI-AMRVAC). In order to limit diffusivity at grid level, we used background and residual decomposition of the magnetic field. The magnetic field is thus made of the sum of a prescribed time-dependent background field B0(t) and a residual field B1(t) computed by the code. In our simulations, B0(t) is essentially made of a rigidly rotating potential dipole field. Results. The first simulation shows that, while plunged in a non-magnetised plasma, a magnetic dipole rotating about an axis oriented at 90° with respect to itself does naturally accelerate the plasma away from the dipole around the rotation axis. The acceleration occurs over a spatial scale of the order of the Alfvénic co-rotation scale r*. During the acceleration, the dipole lines become stretched and twisted. The observed asymptotic fluid velocities are of the order of the phase speed of the fast MHD mode. In two simulations where the surrounding non-magnetised plasma was chosen to move at supersonic speed perpendicularly to the rotation axis (a situation that is reminiscent of Uranus in the solar wind at equinox), the lines of each hemisphere are symmetrically twisted and stretched as before. However, they are also bent by the supersonic flow, thus forming a magnetic tail of interlaced field lines of opposite polarity. Similarly to the case with no wind, the interlaced field lines and the attached plasma are accelerated by the rotation and also by the transfer of kinetic energy flux from the surrounding supersonic flow. The tailwards fluid velocity increases asymptotically towards the externally imposed flow velocity, or wind. In one more simulation, a transverse magnetic field, to both the spin axis and flow direction, was added to the impinging flow so that magnetic reconnection could occur between the dipole anchored field lines and the impinging field lines. No major difference with respect to the no-magnetised flow case is observed, except that the tailwards acceleration occurs in two steps and is slightly more efficient. In order to emphasise the effect of rotation, we only address the case of a fast-rotating planet where the co-rotation scale r* is of the order of the planetary counter-flow magnetopause stand-off distance rm. For Uranus, r*≫ rm and the effects of rotation are only visible at large tailwards distances r ≫ rm.

How a magnetic helix can develop at solstice in a Uranus-type rotating magnetosphere

2020 ◽  
Author(s):  
Léa Griton ◽  
Filippo Pantellini
Keyword(s):  
Rotation Axis ◽  
Double Helix ◽  
Large Angle ◽  
Rotation Period ◽  
Ecliptic Plane ◽  
Characteristic Relaxation Time ◽  
Dipole Axis ◽  
Magnetospheric Tail ◽  
Planetary Magnetic Field ◽  
Fast Rotating

<p class="p1">The characteristic relaxation time of the Uranus magnetosphere is of the order  of the planet's rotation period. This is also the case for Jupiter or Saturn. However, the specificity of Uranus (and to a lesser extent of  Neptune) is that the rotation axis and the magnetic dipole axis are separated by  a large angle (~60°) the consequence of which is the development of a highly dynamic and complex magnetospheric tail. In addition, and contrary to all other planets of the solar system, the rotation axis of Uranus happens to be quasi-parallel to the ecliptic plane which also implies a strong variability of the magnetospheric structure as a function of the season. The magnetosphere of Uranus is thus a particularly challenging case for global plasma simulations, even in the frame of MHD. We present a detailed analysis of MHD simulations of a fast-rotating magnetosphere inspired from Uranus at solstice. At first, a simplified case allows us to explain in detail the formation and the internal structure of a double helix that develops in the magnetotail at solstice. Then we analyse a "real" Uranus simulation with parameters for the solar wind and planetary magnetic field defined from the measurements of Voyager II flyby in 1986.</p>

MHD simulations of an Uranus type rotating magnetosphere

2020 ◽  
Author(s):  
Filippo Pantellini ◽  
Léa Griton
Keyword(s):  
Solar Wind ◽  
Rotation Axis ◽  
Large Angle ◽  
Rotation Period ◽  
Ecliptic Plane ◽  
Characteristic Relaxation Time ◽  
Dipole Axis ◽  
Mhd Simulations ◽  
Planetary Rotation ◽  
Magnetospheric Tail

<p>The characteristic relaxation time of the Uranus magnetosphere is of the order  of the planet's rotation period. This is also the case for Jupiter or Saturn. However, the specificity of Uranus (and to a lesser extent of  Neptune) is that the rotation axis and the magnetic dipole axis are separated by  a large angle (~60°) the consequence of which is the development of a highly dynamic and complex magnetospheric tail. In addition, and contrary to all other planets of the solar system, the rotation axis of Uranus happens to be quasi-parallel to the ecliptic plane which also implies a strong variability of the magnetospheric structure as a function of the season. The magnetosphere of Uranus is thus a particularly challenging case for global plasma simulations, even in the frame of MHD. We present MHD simulations of a Uranus type magnetosphere at both equinox (solar wind is orthogonal to the planetary rotation axis) and solstice (solar wind is orthogonal to the planetary rotation axis) configurations. The main differences between the two configurations will be discussed. </p>

scholarly journals Azimuthal magnetic fields in Saturn’s magnetosphere: effects associated with plasma sub-corotation and the magnetopause-tail current system

2003 ◽  
Vol 21 (8) ◽  
pp. 1709-1722 ◽  
Author(s):  
E. J. Bunce ◽  
S. W. H. Cowley ◽  
J. A. Wild
Keyword(s):  
Magnetic Field ◽  
Solar Wind ◽  
Magnetic Fields ◽  
Current System ◽  
Heavy Ion ◽  
Local Times ◽  
Dynamical Process ◽  
Tail Current ◽  
Field Lines ◽  
Inner Region

Abstract. We calculate the azimuthal magnetic fields expected to be present in Saturn’s magnetosphere associated with two physical effects, and compare them with the fields observed during the flybys of the two Voyager spacecraft. The first effect is associated with the magnetosphere-ionosphere coupling currents which result from the sub-corotation of the magnetospheric plasma. This is calculated from empirical models of the plasma flow and magnetic field based on Voyager data, with the effective Pedersen conductivity of Saturn’s ionosphere being treated as an essentially free parameter. This mechanism results in a ‘lagging’ field configuration at all local times. The second effect is due to the day-night asymmetric confinement of the magnetosphere by the solar wind (i.e. the magnetopause and tail current system), which we have estimated empirically by scaling a model of the Earth’s magnetosphere to Saturn. This effect produces ‘leading’ fields in the dusk magnetosphere, and ‘lagging’ fields at dawn. Our results show that the azimuthal fields observed in the inner regions can be reasonably well accounted for by plasma sub-corotation, given a value of the effective ionospheric Pedersen conductivity of ~ 1–2 mho. This statement applies to field lines mapping to the equator within ~ 8 RS (1 RS is taken to be 60 330 km) of the planet on the dayside inbound passes, where the plasma distribution is dominated by a thin equatorial heavy-ion plasma sheet, and to field lines mapping to the equator within ~ 15 RS on the dawn side outbound passes. The contributions of the magnetopause-tail currents are estimated to be much smaller than the observed fields in these regions. If, however, we assume that the azimuthal fields observed in these regions are not due to sub-corotation but to some other process, then the above effective conductivities define an upper limit, such that values above ~ 2 mho can definitely be ruled out. Outside of this inner region the spacecraft observed both ‘lagging’ and ‘leading’ fields in the post-noon dayside magnetosphere during the inbound passes, with ‘leading’ fields being observed both adjacent to the magnetopause and in the ring current region, and ‘lagging’ fields being observed between. The observed ‘lagging’ fields are consistent in magnitude with the sub-corotation effect with an effective ionospheric conductivity of ~ 1–2 mho, while the ‘leading’ fields are considerably larger than those estimated for the magnetopause-tail currents, and appear to be indicative of the presence of another dynamical process. No ‘leading’ fields were observed outside the inner region on the dawn side outbound passes, with the azimuthal fields first falling below those expected for sub-corotation, before increasing, to exceed these values at radial distances beyond ~ 15–20 RS , where the effect of the magnetopause-tail currents becomes significant. As a by-product, our investigation also indicates that modification and scaling of terrestrial magnetic field models may represent a useful approach to modelling the three-dimensional magnetic field at Saturn.Key words. Magnetospheric physics (current systems; magnetosphere-ionosphere interactions; solar wind-magnetosphere interactions)

scholarly journals Cusp-like plasma in high altitudes: a statistical study of the width and location of the cusp from Magion-4

2002 ◽  
Vol 20 (3) ◽  
pp. 311-320 ◽  
Author(s):  
J. Mĕrka ◽  
J. Šafránková ◽  
Z. Nĕmeček
Keyword(s):  
Magnetic Field ◽  
Solar Wind ◽  
High Altitudes ◽  
Data Set ◽  
Latitudinal Dependence ◽  
Merging Process ◽  
Low Altitude ◽  
Field Lines ◽  
The Magnetic Field ◽  
Dipole Tilt

Abstract. The width of the cusp region is an indicator of the strength of the merging process and the degree of opening of the magnetosphere. During three years, the Magion-4 satellite, as part of the Interball project, has collected a unique data set of cusp-like plasma observations in middle and high altitudes. For a comparison of high- and low-altitude cusp determination, we map our observations of cusp-like plasma along the magnetic field lines down to the Earth’s surface. We use the Tsyganenko and Stern 1996 model of the magnetospheric magnetic field for the mapping, taking actual solar wind and IMF parameters from the Wind observations. The footprint positions show substantial latitudinal dependence on the dipole tilt angle. We fit this dependence with a linear function and subtract this function from observed cusp position. This process allows us to study both statistical width and location of the inspected region as a function of the solar wind and IMF parameters. Our processing of the Magion-4 measurements shows that high-altitude regions occupied by the cusp-like plasma (cusp and cleft) are projected onto a much broader area (in magnetic local time as well as in a latitude) than that determined in low altitudes. The trends of the shift of the cusp position with changes in the IMF direction established by low-altitude observations have been confirmed.Key words. Magnetospheric physics (magnetopause, cusp and boundary layer; solar wind – magnetosphere interactions)

scholarly journals Cellular automata model of magnetospheric-ionospheric coupling

2003 ◽  
Vol 21 (9) ◽  
pp. 1931-1938 ◽  
Author(s):  
B. V. Kozelov ◽  
T. V. Kozelova
Keyword(s):  
Magnetic Field ◽  
Solar Wind ◽  
Cellular Automata ◽  
Interplanetary Magnetic Field ◽  
Control Parameter ◽  
Energy Content ◽  
Nonlinear Phenomena ◽  
Cellular Automata Model ◽  
The Earth ◽  
Magnetospheric Tail

Abstract. We propose a cellular automata model (CAM) to describe the substorm activity of the magnetospheric-ionospheric system. The state of each cell in the model is described by two numbers that correspond to the energy content in a region of the current sheet in the magnetospheric tail and to the conductivity of the ionospheric domain that is magnetically connected with this region. The driving force of the system is supposed to be provided by the solar wind that is convected along the two boundaries of the system. The energy flux inside is ensured by the penetration of the energy from the solar wind into the array of cells (magnetospheric tail) with a finite velocity. The third boundary (near to the Earth) is closed and the fourth boundary is opened, thereby modeling the flux far away from the tail. The energy dissipation in the system is quite similar to other CAM models, when the energy in a particular cell exceeds some pre-defined threshold, and the part of the energy excess is redistributed between the neighbouring cells. The second number attributed to each cell mimics ionospheric conductivity that can allow for a part of the energy to be shed on field-aligned currents. The feedback between "ionosphere" and "magnetospheric tail" is provided by the change in a part of the energy, which is redistributed in the tail when the threshold is surpassed. The control parameter of the model is the z-component of the interplanetary magnetic field (Bz IMF), "frozen" into the solar wind. To study the internal dynamics of the system at the beginning, this control parameter is taken to be constant. The dynamics of the system undergoes several bifurcations, when the constant varies from - 0.6 to - 6.0. The Bz IMF input results in the periodic transients (activation regions) and the inter-transient period decreases with the decrease of Bz. At the same time the onset of activations in the array shifts towards the "Earth". When the modulus of the Bz IMF exceeds some threshold value, the transition takes place from periodic to chaotic dynamics. In the second part of the work we have chosen as the source the real values of the z-component of the interplanetary magnetic field taken from satellite observations. We have shown that in this case the statistical properties of the transients reproduce the characteristic features observed by Lui et al. (2000).Key words. Magnetospheric physics (magnetosphere-ionosphere interactions) – Space plasma physics (nonlinear phenomena)

Theory and observations of switchbacks’ evolution in the solar wind

2021 ◽  
Author(s):  
Anna Tenerani ◽  
Marco Velli ◽  
Lorenzo Matteini
Keyword(s):  
Magnetic Field ◽  
Solar Wind ◽  
Radial Magnetic Field ◽  
Open Questions ◽  
Coronal Activity ◽  
Turbulent Heating ◽  
Field Lines ◽  
Radial Evolution ◽  
Field Correlation

<p>Alfvénic fluctuations represent the dominant contributions to turbulent fluctuations in the solar wind, especially, but not limited to, the fastest streams with velocity of the order of 600-700 km/s. Alfvénic fluctuations can contribute to solar wind heating and acceleration via wave pressure and turbulent heating. Observations show that such fluctuations are characterized by a nearly constant magnetic field amplitude, a condition which remains largely to be understood and that may be an indication of how fluctuations evolve and relax in the expanding solar wind. Interestingly, measurements from Parker Solar Probe have shown the ubiquitous and persistent presence of the so-called switchbacks. These are magnetic field lines which are strongly perturbed to the point that they produce local inversions of the radial magnetic field. The corresponding signature of switchbacks in the velocity field is that of local enhancements in the radial speed (or jets) that display the typical velocity-magnetic field correlation that characterizes Alfvén waves propagating away from the Sun. While there is not yet a general consensus on what is the origin of switchbacks and their connection to coronal activity, a first necessary step to answer these important questions is to understand how they evolve and how long they can persist in the solar wind. Here we investigate the evolution of switchbacks. We address how their evolution is affected by parametric instabilities and the possible role of expansion, by comparing models with the observed radial evolution of the fluctuations’ amplitude. We finally discuss what are the implications of our results for models of switchback generation and related open questions.</p>

scholarly journals Instabilities of Magnetic Fields in Stars

1974 ◽  
Vol 59 ◽  
pp. 177-177
Author(s):  
R. J. Tayler
Keyword(s):  
Magnetic Field ◽  
Gravitational Potential ◽  
Rotation Period ◽  
Region Of Interest ◽  
Alfven Wave ◽  
Closed Field ◽  
Toroidal Plasma ◽  
Wide Range ◽  
Field Lines ◽  
Rotating Star

It has been shown (Markey and Tayler, 1973; Tayler, 1973; Wright, 1973) that a wide range of simple magnetic field configurations in stars are unstable. Although the ultimate effect of the instabilities is unclear, it seems likely that they would lead to enhanced destruction of magnetic flux, so that magnetic field decay would be much more rapid than previously supposed. Instability is almost certain in a non-rotating star containing either a purely toroidal field or a purely poloidal field, which has closed field lines inside the star. In both cases the instability resembles the well known instabilities of cylindrical and toroidal current channels, modified by the constraint that motion must be almost entirely along surfaces of constant gravitational potential.If both toroidal and poloidal fields are present, the problem is more complicated. In a toroidal plasma with a helical field, the worst instabilities are also helical but it is impossible for a helical disturbance to be parallel to a surface of constant gravitational potential everywhere. As a result, the admixture of toroidal and poloidal fields has a stabilizing influence, but it is not at present clear whether the majority of such configurations are completely stable.The effect of rotation has not yet been studied but it will certainly be important if the rotation period is less than the time taken for an Alfvén wave to cross the region of interest. This is true in most stars unless the internal magnetic field is very much stronger than any observed field.

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