The Sun and Cool Stars: activity, magnetism, dynamos

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An ICME is an interplanetary expansion of a coronal mass ejection CME , where a large coronal loop breaks down, emitting vast amounts of solar material into the interplanetary space. The CMEs are related to magnetically active solar regions and reflect the solar toroidal phase, having maximum occurrence around sunspot maxima.

On the other hand, the HSSs arise from large empty regions of solar corona called the coronal holes CH , which normally develop around the solar poles during the declining phase of the solar cycle. While sunspots and toroidal fields are decreasing at this time, the intensity of solar polar fields is increasing and the solar magnetic field approaches the poloidal phase.

However, this development is not uniform and often there are long extensions of polar coronal holes toward the solar equator. Such longitudinally asymmetric CH structures, frequently called "elephant trunks", are important because they emit fast solar wind over a wide latitude range. When the fast HSS stream from a trunk reaches the prior slow solar wind, an important interaction region arises with favorable conditions for magnetic storms and auroral substorms.

These regions are called corotating interaction regions CIRs because they can repeat several times in successive solar rotations. These results suggest that the solar poloidal phase, despite its superficially inactive nature, is at least as important for space weather and other consequences of solar activity in the near-Earth environment as the solar toroidal phase. This further underlines the need of knowing the development of the poloidal phase more accurately and over a longer time interval than so far.

Magnetism is ubiquituous in the universe - planets, stars, galaxies and even the intergalactic medium are all magnetized. The modern society is strongly affected by the magnetic activity manifestations of our star, the Sun.

The space missions, air traffic, power grids and telecommunication networks are vulnerable to disturbances caused by bad space weather, i. The magnetized objects share a few common features: at least they are rotating and the matter in which the magnetic fields are generated interior of the Earth, solar convection zone, interstellar matter is turbulent. Understanding the generation of magnetic fields, therefore, is intrinsically linked to studying turbulent fluids under extreme conditions.

Such studies are out of the scope of analytic methods and laboratory experiments, the only viable tool being numerical modelling. Even that approach is extremely challenging: to be able to understand, explain, and predict the behavior of the global-scale magnetic field of the Sun, one needs to build a numerical model of the fine details of the motions within the star.

This can be achieved only through by developing the fastest and the most accurate possible algorithms that are suitable for massively parallel computation, and continuously improving them to take advantage of the new emerging ICT technologies such as cloud and GPU computing.

As a result of the modelling, huge datasets with three-dimensional spatial information of the basic physical quantities such as velocity, magnetic field, density, temperature are obtained. As the magnetohydrodynamic equations are integrated over time, a statistically steady state is usually obtained, but the solutions most often show systematic behavior large-scale circulation, magnetic cycles making it necessary to study also the time dimension of the data.

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In this way, even the moderated-sized simulations create big data when all the information is build into a spatiotemporal stack. With increasing amount of data, the development of efficient and computationally economical data analysis methods is an urgent requirement in our research group. The lines of study described above constitute a major part of a discipline called computational astrophysics. In the year , computational astrophysics user group, of which roughly half of the projects are lead by the CMDAA group researchers, consumed roughly core hours per user of the national computing centre CSC resources.

This is the top number amongst all the disciplines, and reflects the large volume of computational resources required for the execution of the models. Systematic observations of the Sun have been carried out for over four hundred years, during the last century with increasing amount of space-borne machinery. Complemented with the data obtained for other solar-type stars coming from many observational infrastructures scattered on the best observing sites over the globe and some satellites, the amount of data is accumulating with increasing speed also in the observational frontier.

Investigation of these objects is important, as this sheds some light into the history of the Sun, as most of the targets observable to Earthly astronomers are rapid rotators, resembling the Sun at an earlier age. The most common method to search for systematic cyclic behavior as observed on the Sun , indicative of a magnetic activity cycle, is to use time-series analysis methods on stellar photometric data. For astronomical observations in the optical range of wavelengths made from the ground, the observations are inevitably interrupted during daytime and due to bad weather.

In astronomical community, therefore, a special class of time series analysis methods capable of effective functioning even in the presence of long time intervals of missing data called as gaps has been developed and are extensively applied in the stellar community.

Cool Stars Have Magnetic Fields Too

All of the equilibria constructed prove to be unstable, and the authors tentatively suggest that stable equilibria might not exist in barotropic stars. In the light of these results, it is probably safe to assume that barotropic stars cannot host MHD equilibria.

Cool Experiment Sand, Magnet & Iron filings

As neutron stars do host magnetic fields, it seems they cannot be perfectly barotropic, or that the crust plays an important role. This general principle is seen in various contexts in MHD; an early reference is Chandrasekhar [ ], sections 84 and after, and fig.

The idea of the Coriolis force, rather than inertia, balancing whatever is driving fluid motion, is of course also well known from atmospheric physics quasi-geostrophic balance, e. It is easy to reconcile the predictions with observed field strengths of 0. Finally, if it can be assumed in all stars which hosted a pre-main-sequence convective dynamo, that the dynamo leaves behind a magnetic field, then a magnetic field of this order of magnitude should be visible at the surface of all such stars during the main sequence.

However, this assumption is not certain—the slow retreat of the pre-main-sequence convective envelope containing a time-dependent dynamo will leave a field of small radial length scale, causing the field to decay more quickly via magnetic diffusion. Its survival will also be affected by processes like meridional circulation and differential rotation. Various explanations spring to mind, the most obvious of which is Ohmic diffusion.

The global timescale for Ohmic diffusion is of the order of 10 10 years, somewhat longer than the main-sequence lifetime of the least massive stars in question. In the absence of other effects, one would expect the electric current associated with the magnetic field to die away reasonably quickly in the outer part of the star, so that after some time the field at the surface is simply a potential-field extrapolation of the field further inside. Depending on the initial geometry and radial distribution of the magnetic field, this could cause the surface field either to rise or fall during the main sequence.

This will depend on the origin of the magnetic field. The main weakness of finite conductivity in explaining this decay though is that it is also observed in massive stars with much shorter main-sequence lifetimes and somewhat longer Ohmic timescales than intermediate-mass stars. Another possibility is the combination of buoyancy and thermal diffusion. In short, a magnetic feature is in pressure balance with its less-strongly magnetized surroundings, and so its gas pressure must be lower; to avoid moving on a dynamic timescale its temperature must therefore be lower than its surroundings.

Heat consequently diffuses into it, resulting in a buoyant rise to the surface [ 65 , 67 , 85 , 93 , ]. Note that this mechanism is distinct from the so-called buoyancy instability or Parker instability where diffusion is not required.

1. Introduction

On the other hand, it might be tricky to get this process to work fast enough: even assuming that the interior field is 10 times stronger than the surface field, the timescale for the most strongly magnetized stars e. However, as with Ohmic diffusion, it may be possible to massage the numbers in the light of the fact that the thermal diffusion timescale is much shorter near the surface of the star.

Another possibility is meridional circulation. The geometry of this flow has been clarified in the seminal work of Zahn [ ].

IAU Colloquium The Sun and Cool Stars, Activity, Magnetism, Dynamos - CERN Document Server

Note the similarity with the timescale of diffusive buoyancy in the previous paragraph—the only difference is that rotational energy has replaced magnetic energy. As the ratio of energies in this case can be much closer to unity than in the previous case, the timescale can also be much smaller, indeed the fast-rotating magnetic stars should experience relatively fast decay, unless the magnetic field finds some way to coexist alongside the meridional flow.

It may be, though, that this flow is simply inhibited by the magnetic stress in the fossil field. Essentially nothing is present in the literature regarding how this inhibition might work. Alternatively, an interaction with the convective core could be the crucial process—the convective motion and waves that it sends into the radiative envelope could somehow result in an enhanced diffusion.

In this case, one would certainly expect a correlation with mass, as the core is very small in late A stars but reaches to around a third of the stellar radius in O stars. The discussion in the previous section begs a question: where does the variation in magnetic fields in otherwise similar stars come from?

According to the traditional model, variations in magnetic field strength in the interstellar medium ISM are simply carried forwards into the star. In the light of recent results though, including the very weak fields in Vega and Sirius, this scenario, at least in its simplest form, now looks very unlikely—the range in field strengths in stars is far greater than that in the ISM.

In addition, this model requires an additional ingredient to produce the observed bimodality between Ap and other A stars. Also, this model is compatible with the lack of magnetic stars observed in binaries, because collapsing cloud cores with a strong magnetic field will spin down efficiently, whereas cores lacking a strong field will retain too much angular momentum to form a single star and so form a binary.

This effect has been seen in simulations [ ]. To summarize, the simple ISM-variation model ignores much of the star formation process and so, alone, it will not explain what we see in stars.

It may play some role however; in any case it is worthwhile to take a more detailed look at star formation from the perspective of magnetic field evolution. Once a cloud has become supercritical, it can collapse dynamically. There is additional evidence for the presence of ordered magnetic fields in discs.

Systems ranging from protostars to active galactic nuclei usually not always, and not all of them show evidence of a fast outflow in the form of a collimated jet. The default model for its origin is the rotation of an ordered magnetic field in the inner regions of the disc. That is, a field crossing the disc with a uniform polarity over a significant region around the central object. Models assuming the existence of such an ordered field as opposed to the small-scale field of mixed polarities generated in magnetorotational instability MRI turbulence have been particularly successful in producing fast magnetically driven outflows.

Accepting this as evidence for the existence of such ordered fields, they might also be the fields that are accreted to form magnetic A, B and O stars. The origin of the ordered field in discs is less certain. It can change only by field lines entering or leaving the disc through its outer boundary, i. In other words, we have an extra phenomenon to explain: why even the most strongly magnetized stars have such weak fields.

This may be related somehow to the fact that accretion discs are turbulent. Simple estimates show that accretion of an external field is very inefficient if the disc has a magnetic diffusivity similar to the turbulent viscosity that enables the accretion [ ].

Indirect imaging of stellar surfaces

Numerical simulations [ ] show that this is in fact a good approximation for magnetorotational turbulence. Though intuitively appealing, accretion of the field of a protostellar cloud as the source of Ap star fields is therefore not an obvious possibility. The flux bundles that drive jets from the inner regions of the disc, inferred indirectly from observations, must somehow be due to a more subtle process.

The Sun and Cool Stars: activity, magnetism, dynamos The Sun and Cool Stars: activity, magnetism, dynamos
The Sun and Cool Stars: activity, magnetism, dynamos The Sun and Cool Stars: activity, magnetism, dynamos
The Sun and Cool Stars: activity, magnetism, dynamos The Sun and Cool Stars: activity, magnetism, dynamos
The Sun and Cool Stars: activity, magnetism, dynamos The Sun and Cool Stars: activity, magnetism, dynamos
The Sun and Cool Stars: activity, magnetism, dynamos The Sun and Cool Stars: activity, magnetism, dynamos

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