Thursday, 19 December 2013

Solar Facts I

Bands for spectroscopy/imaging
Band Wavelength (nm) Region
G 430.55 photosphere
FeI 630.2 photosphere/low chromosphere
H-alpha 656.28 Chromosphere
CaII 854.00 Chromosphere
HeII 30.4 Chromosphere
FeIX 17.1 Corona/transition region

CaII spans a range of atmospheric heights. From the far wings into the core the line samples from the low to high photosphere. The line core itself is formed in the lower chromosphere.

The beta of a plasma, symbolized by β, is the ratio of the plasma pressure (p = nkBT) to the magnetic pressure (pmag = B²/2μ0).
For plasma with β << 1 the pressure may be neglected.
For plasma with β >> 1 the Lorentz force may be neglected.

The plasma β for different plasma is shown in the table below
Plasma type Coronal Loop Active Corona Coronal Hole Solar wind Magnetosphere
B 300G (0.03T) 100G (0.01T) 10G (0.001T) 6x10-9 T 3x10-5 T
beta 0.0004 0.01 0.007 1 4x10-6
T (K) 106 2x106 106 105 104
Density (kg/m3) 1.7x10-11 1.67x10-13 1.67x10-13 1.7x10-20 1.7x10-17
Alfven Speed(km -1 6500 41 6500
Sound Speed(km -1 120 37 12
Magnetic Reynolds number 6.8x1013 1.5x1012 1.7x10-17

Saturday, 7 December 2013

Dynamics of the Solar Atmosphere Generated by the Eigen Modes of Solar Global Oscillations

This post describes progress with a study Dynamics of the Solar Atmosphere Generated by the Eigen Modes of Solar Global Oscillations.


The solar atmosphere exhibits a diverse range of wave phenomena one of the earliest to be discovered was the five minute oscillation, the p-mode. The solar p modes are generated by global resonant oscillations and turbulent motions just beneath the photosphere. The resulting propagation of this wave energy into the solar atmosphere may be used as a diagnostic tool to predict some of the physical characteristics of the  suns atmospheric layers.Report a study of synthetic photospheric oscillations in both non-magnetic and magnetic  model of the quiet sun


To investigate the dynamics in the solar atmosphere which are generated by solar global eigenmodes of oscillation and to understand mechanisms of leakage of 5 min global oscillations into the atmosphere. Understand the conditions under which chromospheric dynamics evolve as a result of the 5 minute global oscillations - (spicules, waves)


This post presents a series of hydrodynamic simulations modelling a realistic solar atmosphere using a driver located at 0.5Mm above the temperature minimum. A combination of the VALIIIC and McWhirter solar atmospheres and coronal density profiles were used as the background equilibrium model in the simulations. With the objective of recreating atmospheric motions generated by global resonant oscillation the driver is spatially structured and is extended in a sinusoidal profile arcoss the base of the computational model. Vertical and horizontal harmonic sources, located at the footpoint region of the open magnetic flux tube, are incorporated in the calculations, to excite oscillations in the domain of interest. To carry out the simulations, we employed the MHD code SMAUG (Sheffield MHD Accelerated Using GPUs). This code is a version of the Sheffield advance code implemented on Graphical processing units [Shelyag2008].


Our results demonstrate the transition of modes with a period of 30s to a mode with a period of 180s,. This hydrodynamic effect demonstrating that the Chromosphere is a source of 180s modes. 

The simulations described were performed using a grid of size 128x128x128. The box size in the vertical direction (the z-direction corresponding to the direction of the solar gravitational field) was 6Mm. The box dimensions in the x and y direction was 4Mm. Continuous boundary conditions were applied to all sides of the box.

Above is a movie of the evolution of Vz showing the development of the initial perturbation in the nonmagnetic equilibrium generated by the 30-second-period driver (in ms−1). The z-axis corresponds to height measured in megameters and the x and y horizontal axes are parallel to the solar surface. 

Above is a distance time plot for the fundamental mode with 30s period for the z and y component of the velocity. The section  is taken at 0.94Mm across the box.

Three-dimensional movie of the evolution of Vz showing the development of the initial perturbation in the nonmagnetic equilibrium generated by the 180-second-period driver (in ms−1). The z-axis corresponds to height measured in megameters and the x and y horizontal axes are parallel to the solar surface. 

Above is a distance time plot for the fundamental mode with 180s period for the z and y component of the velocity. The section  is taken at 0.94Mm across the box.

Three-dimensional movie of the evolution of Vz showing the development of the initial perturbation in the nonmagnetic equilibrium generated by the 300-second-period driver (in ms−1). The z-axis corresponds to height measured in megameters and the x and y horizontal axes are parallel to the solar surface. 

Above is a distance time plot for the fundamental mode with 300s period for the z and y component of the velocity. The section  is taken at 0.94Mm across the box.

The results presented below are for the same driver with the period of 30s and with a magnetic field of 10G. Three different magnetic configurations are considered.

  • Fixed vertical field of 10G
  • Flux tube of width 1Mm located at the centre of the model and maximum field magnitude of 10G
  • Flux tube of width 2Mm located at the centre of the model and maximum field magnitude of 10G

The flux tube field configurations were computed using the, self consistent field method (e.g. see notes from Fred Gents presentation).

Vertical 10G Field

The diagram above shows the magnetic Field Configuration, Alfven Speed, Sound Speed and Temperature Profile for model with constant vertical 10G Magnetic Field.

Above is a movie of the evolution of Vz showing the development of the initial perturbation in the nonmagnetic equilibrium generated by the 30-second-period driver (in ms−1) and with the vertical 10G magnetic field. The z-axis corresponds to height measured in megameters and the x and y horizontal axes are parallel to the solar surface. 

1Mm 10G Flux Tube - Using self similar construction

The diagram above shows the magnetic Field Configuration, Alfven Speed, Sound Speed and Temperature Profile for model with 1Mm width Flux Tube.

Above is a movie of the evolution of Vz showing the development of the initial perturbation in the nonmagnetic equilibrium generated by the 30-second-period driver (in ms−1) and with the 1Mm magnetic flux tube of maximum strength 10G. The z-axis corresponds to height measured in megameters and the x and y horizontal axes are parallel to the solar surface. 

2Mm 10G flux tube using self similar field construction

The diagram above shows the magnetic Field Configuration, Alfven Speed, Sound Speed and Temperature Profile for model with 2Mm width Flux Tube.

Above is a movie of the evolution of Vz showing the development of the initial perturbation in the nonmagnetic equilibrium generated by the 30-second-period driver (in ms−1) and with the 2Mm magnetic flux tube of maximum strength 10G. The z-axis corresponds to height measured in megameters and the x and y horizontal axes are parallel to the solar surface. 

Friday, 22 November 2013

Energy Loss of Solar f- and p-Modes due to MHD tube wave excitation.

I'm currently running numerical MHD simulations of the dynamics in the solar corona generated by Solar Global oscillations.

Here are some key references I'm studying

  • Helioseismology by J. Christensen-Dalsgaard
  • Lecture Notes on Stellar Oscillations by J

  • Today there was an interesting seminar entitled "Energy Loss of Solar f- and p-Modes due to MHD tube wave excitation" as it seems fairly close to some of the areas I'm investigating this seemed to be a must attend event! The talk was a discussion of the generation and propagation of sausage tube waves within the solar convection zone and chromosphere by the buffeting of p modes. Evidence of p-modes and solar global oscillations is nowadays easily observed from satellite imagery. The images below from SDO for the 15th November 2013 are the Magnetogram, HMI Intensity gram and Dopplergram.


    The Magnetogram shows the magnetic strengths and opposite polarities in the vicinity of sunspots. The right hand image below is a dopplergram. It is interesting to note that the dopplergram shows a suppression of the surface velocity at the at sunspot locations and suggesting that the intensity of p modes are reduced by magnetic fields. The observed power spectrum for p-mode oscillations is shown in the diagram below the intense region at 3.3mHz corresponds to the ubiquitous 5 minute oscillations. It is important to note the ridges present in this diagram. 

    The period versus horizontal wavelength diagram obtained by the Michelson Doppler Imager (MDI) on the SOHO spacecraft. Since only waves with specific combinations (related to the Sun's interior structure) of period and horizontal wavelength resonate within the Sun, they produce the fine-tuned `ridges' of greater power. (Image source: SOHO). The power spectrum obtained from  the MDI Medium-l data for the modes averaged over the azimuthal order m. The power concentrates in ridges corresponding to solar acoustic (p) modes. The lowest weak ridge corresponds to the fundamental (f) mode.

    Simulations may be used to excite a rich spectrum of p-mode oscillations (left), very similar to the MDI diagram (right). The dark line is the theoretical f-mode.

    The power spectrum above represents p-mode frequency oscillations of the Sun. The x and y axes are wave numbers which are calculated from the degree l of the spherical surface harmonic modes. The z axis is the frequency of oscillation.  In the diagram above the rings cut out of the z-axis plane are 3D representations of the p-mode ridges. Six rings equal six modes. Eight rings equal eight p-modes, and so on. The frequency value of the top of the z plane is the Nyquist frequency, the high-frequency limit given by the time resolution of the signal, which here is 8334 microHz. The densest part of the ridge pattern is roughly 3000 microHz. The frequency resolution is 5.71 microHz. This data was acquired using the Michelson Doppler Imager instrument on board the SOHO spacecraft.

    For solar oscillations, the power is not evenly distributed in the k-omega plane, but instead corresponds to ridges. Each of these ridges corresponds to a fixed number of wave nodes in the radial direction. The ridges are seen here as "rays" in the cut planes facing you.

    Introduction to helio-seismology

    Tube waves propagate along the many magnetic fibrils which are embedded in the convection zone and expand into the chromosphere due to the fall in density with height of the surrounding plasma.   The magnetic fibrils form a waveguide for these waves to freely propagate up and down the tube, those waves propagating upward pass through the photosphere into the chromosphere and enter the upper atmosphere, where they can be measured as loop oscillations and other forms of  propagating coronal waves. We treat the magnetic fibrils as vertically aligned, thin flux tubes embedded in a two region polytropic-isothermal atmosphere to investigate the coupling of p-mode driven sausage waves; which are excited in the convection zone and propagate into the overlying chromosphere. 

     Observational evidence for tube modes

    Observational evidence of tube modes presented by Srivistava et al (ApJ 2013)

    The excited tube waves carry energy away from the p-mode cavity resulting in a deficit of p-mode energy which we quantify by computing the associated damping rate and absorption coefficient of the driving p modes.  We calculate the damping rates/absorption coefficients and compare them with observations and previous theoretical studies of this nature.

    It is understood that magnetic field lines suppress the p-mode oscillations this occurs as a result of sausage mode oscillations being driven by the p-modes. The mode conversion from acoustic to magneto acoustic tube mode results in a substantial absorption of energy

    Friday, 8 November 2013

    Solar atmospheric MHD flux tube equilibria

    Fred Gent's presentation reviewed an area which is vital for correct modelling solar phenomena driven by p-mode oscillations. Numerical modelling of wave phenomena in the solar atmosphere require information about the structure of magnetic fields therein. Such information is gleaned from magnetograms or from the analysis of Zeeman splittings or Stokes profiles. Improvements are expected as coronal seismology increases in maturity and can be used to infer information about the magnetic field. Even when armed with this range of sophisticated spectropolarimetric methods it is challenging to construct a numerical model of the magnetic fields pervading the solar atmosphere and which can be used for simulations of solar wave phenomena. For an example see Shelyag et al 2010.

    The seminar described an analytical method for constructing the magnetic field using the magnetohydrostatic balance equation. Currently we use a numerical approach based on the self-similarity assumption ensuring that the divergence of the constructed magnetic field is zero everywhere. This construction method was considered by Schlüter, A., & Temesváry, S. 1958 and Schüssler & Rempel 2005. The construction is shown below.

    The objective is to construct a 3D magnetic flux tube in magnetohydrostatic equilibrium within a realistic stratified solar atmosphere. The ambition is to extend this to model multiple flux tubes in quasi-equilibrium. The magnetic configuration is comprised of vertical flux tubes expanding with height in response to the fall in plasma pressure, and the adjustment of the plasma pressure and density distributions arise from the analytic solution of the pressure balance equation.

    The plot above shows interpolated 1D fits to vertical hydrostatic atmospheric profiles (Vernazza et al. 1981; McWhirter et al. 1975, former up to 2.3 Mm, latter above 2.4 Mm): thermal pressure p (Pa) (dotted, light blue to blue), plasma density ρ ( μg m− 3) (dashed, purple to yellow) and temperature T ( K) (dash–dotted, red to green). Using these profiles it is possible to construct density and pressure profiles for a solar atmosphere in quasi-hydrostatic equlibrium. The profiles shown below are initially computed for a field free solar atmosphere the reference values for density and pressure, these are obtained from the measured data.

    The next stage in the process is to consider the contribution from the magnetic pressure
    Using a modified form of the fields for the self similarity method, it is possible to explicitly integrate the expressions for the magnetostatic balance. A single open magnetic flux tube spanning the solar photosphere (solar radius ≃ R) and the lower corona (R + 10 Mm) was modelled in magnetohydrostatic equilibrium within a realistic stratified atmosphere subject to solar gravity.  The results of the analysis for an axially symmetric 3D structure, with magnetic field strength, plasma density, pressure and temperature all consistent with observational and theoretical estimates is shown in the diagrams below;

    The left-hand diagram above shows a 3D rendition of the magnetic flux tube including the magnetic field lines (reducing field strength, turquoise–blue). The rear and bottom surfaces display the thermal pressure (reducing, brown–yellow) and the isosurfaces depict plasma-β (purple–green ≃277,1,0.08,0.025and0.016). A vertical 2D slice of the magnetohydrostatic background magnetic pressure is illustrated in the middle image. Some representative field lines are overplotted in blue. The box (black, dotted) encloses the region magnified for display in the image on the right.

    Solar Flux tubes  are observed to remain relatively stable for up to a day or more, and one of the objectives here is to apply the model as the background condition for numerical studies of energy transport mechanisms from the surface to the corona. 

    The diagram above shows a vertical 2D slice log profile of the magnetohydrostatic background (a) thermal pressure p (b) density ρ and (c) temperature T. Magnetic field lines (solid, blue) are overplotted in (a) and (b). The diagram below shows a vertical 2D slice of the log magnetohydrostatic background plasma-β: the ratio of thermal to magnetic pressure. 

    The self-similar construction ensures the magnetic field is divergence free. The equation of pressure balance for this particular set of flux tubes can be integrated analytically to find the pressure and density corrections required to preserve the magnetohydrostatic equilibrium. The model includes a number of free parameters, which makes the solution applicable to a variety of other physical problems and it may therefore be of more general interest. The presentation generated a lot of discussion one of the questions was around  the use of a consistency rule which implied a current free region. What was particularly exciting was the presentation of magnetic field configurations featuring multiple flux tubes 

    Ref: F.Gent et Al Monthly Notices of the Royal Astronomical Society, Volume 435, Issue 1, p.689-697
    (full journal article here)

    (arxIv version )

    Tuesday, 15 October 2013

    Cycle to Cycle Variations in Solar Activity

     One of the first seminars for the Solar Wave Theory Group this year was a highly informative presentation entitled "Cycle to cycle variations in solar activity: oscillator models and prospects for cycles 24 and 25". The first part  of this seminar was a presentation of a systematic study of the parameter space of the van der Pol and van der Pol-Duffing oscillators, identifying parameter domains where a behaviour analogous to the observed characteristics of the sunspot number series is shown. The implications of the findings for the solar dynamo were discussed. In the second part we were  presented with a discussion of the current status of solar activity and forecasts for cycles 24 and 25.

    Galileo discovered sunspots around 1612, since then, ongoing study has revealed an 11 year cycle. This seminar lived upto its expectations. A four hundred year record of Sunspot observations is shown below. The Maunder minimum is the long period shown very clearly from 1650-1700 whilst the Dalton Minima occurred between 1800 and 1825.

    In order to understand these phenomena we understand the solar magnetic field as a dipole at high polar latitudes. It is believed that the variations constituting the solar cycle arise from the hydro-magnetic dynamo which starts at the bottom of the convective zone of the sun. Tubes of magnetic flux formed in the solar interior and resulting from turbulent mixing  it takes around 5-6 years (half of the solar cycle time) for these to appear as sunspots on the solar surface. The duration of the solar cycle and the distribution of sunspots on the solar surface is therefore determined by the mechanisms which affect the travel time of flux tubes to the solar surface. The sunspot distribution visualised using the Butterfly diagrams (for example see below) results from the motion in longitude and latitude of the sunspots in the solar surface. It can be seen that a minority of sunspots migrate towards the solar poles whilst the majority goto the equator.  

    These effects are determined by the nature of the solar dynamo the diagrams below illustrate how the field distribution is determined by the differential rotation. The so called flux transport dynamo is used by some models to simulate and predict solar cycles.

    The Hale law describes the nature of the butterfly diagram. Sunspots form in pairs and the polarity of the leading spot matches the polarity of the polar cap (in the same hemisphere), this is called Hales polarity law. The line connecting the leading spot to the following spot is tilted toward the equator, the tilt angle increases with solar latitude, this is called Joy's Law. The Waldmeier effect is the relation that the rise time in months is approximately 35 + 1800 /(amplitude in sunspot number). The “Waldmeier effect,” is an anticorrelation between the peak in sunspot number of a cycle and the time from minimum to reach that peak. These effects can be attributed to the solar dynamo and good simulations will predict these distribution laws.

    We were presented with an interesting table comparing the predictive capabilities of different magneto-dynamo models

    The final part of the seminar was fascinating as it presented the interesting idea of the reduction of sunspot numbers and the possibility of a grand minima (e.g. the Maunder minimum). This is currently controversial as it has been suggested that this deepest of solar minima may result in the disappearance of sunspots for cycle 25!


    Tuesday, 17 September 2013

    First version of CUDA Enabled SAC Code: SMAUG

    The first version of the CUDA enabled SAC code, called SMAUG is available for download and testing in this post we provide details for downloading the current distribution.

    The Sheffield Advanced Code (SAC) is a novel fully non-linear MHD code, designed for simulations of linear and non-linear wave propagation in
    gravitationally strongly stratified magnetised plasma.

    The GPU version of the code may be downloaded from

    Getting started documentation for the code is available at SMAUG Requires the following hardware and software

    CUDA-Enabled Tesla GPU Computing Product with at least compute capability 1.3.

    CUDA toolkit

    Correctly installed and compiler on user path.

    Benchmarking SAC on Intel and AMD Processors

    This blog entry presents tables of timings for running the SAC code on systems with different intel and AMD cpu's. Test code was provided by Viktor Fedun and is a SAC model of a flux tube. The work was undertaken in January 2011. The test problem was allowed to run  for the first 20 iterations and the totsl time recorded. Theproblem was also run on a diffrent number of processing cores.

    SWAT server Intel Nehalem E5530 @2.4GHz
    OpenMPI compiled with gigabit ethernet
    Num. ProcsTime/step(sec)Total time(sec)
    2 12.3 246.09
    4 9.76 195.28
    6 9.36 187.21
    8 10.98 175.65
    10 19.758 197.58
    12 15.53 186.45
    14 17.7 248.1
    16 13.5 216.3

    ICEBERG AMD Opteron Barcelona 2376 2.3GHz
    OpenMPI Using Infiniband
    Num. ProcsTime/step(sec)Total time(sec)
    16 6.39 127.71
    14 6.49 129.78
    12 8.13 162.5
    10 7.73 154.69
    8 8.51 170.16
    6 10.13 202.64
    4 12.66 253.1
    2 22.63 452.7

    ICEBERG AMD Opetron Barcelona 2376 2.3GHz

    OpenMPI gigabit ethernet

    Num. ProcsTime/step(sec)Total time(sec)
    2 23.0 460.0
    4 13.5 269.6
    6 11.3 225.9
    8 11.0 220.6
    10 8.0 160.2
    12 10.3 206.5
    14 7.6 152.7
    16 6.5 130.0

    ICEBERG DELL C60100 testnode AMD INTEL Westmere EP (Gulftown 6c) Xeon X5650 2.67GHz
    OpenMPI gigabit ethernet
    Num. ProcsTime/step(sec)Total time(sec)
    2 11.2 224.6
    4 7.9 157.8
    6 7.3 146.1
    8 6.6 131.2
    10 5.9 118.0
    12 6.1 121.0
    14 5.7 113.6
    16 5.0 100.3
    18 4.7 93.9
    20 4.6 92.8
    24 4.1 81.3

    ICEBERG AMD Opteron Barcelona 2347 1.9GHz
    OpenMPI Using Infiniband
    Num. ProcsTime/step(sec)Total time(sec)
    16 7.18 114.88



    8 9.77 195.43


    2 23.23 464.67

    vac3d test AMD Opetron6176  12cores and 2300MHz using the gateway test cluster, MPICH . Qlogic infiniband

    Num. ProcsTime/step(sec)Total Time(sec)

    vac3d test AMD Opteron6140  8cores and 2600MHz using the gateway test cluster, MPICH . Qlogic infiniband

    Num. ProcsTime for one step(s)Total time (sec)

    vac3d ICEBERG DELL C60100 newworkers AMD INTEL Westmere EP (Gulftown 6c) Xeon X5650 2.67GHz using /fastdata
    Num. ProcsTime for one step(s)Total time (sec)
    1 15.30 305.936
    2 10.40 207.956
    4 5.50 109.953
    6 3.76 75.28
    8 4.56 91.23
    9 3.69 73.72
    10 3.22 64.301
    12 2.74 54.80
    14 2.49 49.72
    16 2.31 46.17
    32 1.51 30.20
    64 1.13 22.57
    125 1.26 25.21
    216 1.06 21.25
    384 0.94 18.77

    idl6.2 gives Segmentation fault under ubuntu

    When executing routines such as tvframe using idl6.2 under ubuntu gives a segmentation fault.

    The tip from fanning consultation for resolving this can be found here

    One resolution is to down load and unpackage the libX11-6_1.0.3-6_i386.deb file.

    This file can be downloaded from

    From the directory from which you will run idl extract the .deb file
    using the command

    dpkg-deb -x libx11-6_1.0.3-6_i386.deb

    Next add the following line to the .bashrc file

    export LD_PRELOAD_PATH=/my/private/lib

    for tcsh (using .tcshrc )

    setenv LD_PRELOAD_PATH /my/private/lib

    Here is a quote from another forum which makes the point

    "this my/lib.... has to changed with the path where you put old file. For example I am using fedora7 with bash and I put old libX11 file under /home/orek then I add export LD_PRELOAD_PATH=/home/orek line in my .bashrc file.
    After modifying .bashrc open new terminal window and run seadas. The trick here is run seadas under the same path, like if you put the libX11 under /home/orek you should run seadas from /home/orek then you may change the working from seadas easily if you like. If you do not have old file I can send you via email it is 1 mb file. For editing .bashrc file you may use gedit or other text editors present in Linux.."

    Saturday, 8 June 2013

    MHD Waves in the Chromosphere

    We are continuing the development of the SMAUG, a GPU based code for computational MHD, the code will be used to study wave phenomena in the solar atmosphere. 

    I attended an interesting presentation by Richard Morton reviewing wave phenomena in the Chromosphere. The Chromosphere is a layer in the sun's atmosphere which is about 2000km thick and is situated above the photosphere.  The Hydrogen in the Chromosphere is partially ionised, features many flows, waves and is highly dynamical. Through the comparatively transparent Chromosphere we observe granulation in the photosphere and features such as sunspots, pores at higher resolution we can observe the inter-granular brightpoint features.

    Studies of the Chromosphere are undertaken using spectral analysis of the core of the H-alpha line[Rutten2008]. Magnetic features can be understood using G band studies and the wings of the H alpha line. Leenaarts et al 2012 undertook 3D modelling of the Chromosphere to achieve realistic line formation. Studies of the CaII line have been used to demonstrate that Fibrils may trace magnetic structures in the Chromosphere [de la Cruz Rodríguez 2011]. Other features include the ubiquitous short lived dynamic eruptions observed on the solar limb, known as spicules.

    Wave motions in the solar atmosphere may be used for
    • understanding mechanisms for heating the solar corona
    • magneto-seismological studies for measuring the state of the solar atmosphere

    There are numerous possibilities for wave mode propagation in the solar environment. An example of a p-mode waveform is shown in the following video, generated using SMAUG, this illustrates the leakage of energy into the solar atmosphere. This energy results from a pressure wave driver in the photosphere. In the case below the driver has a period of 30s and drives the lower boundary of this 2D model with half of a sinusoidal wave form.

    Further examples of 3D computational models of the leakage of photospheric motions in the solar atmosphere may be found at Fedun2009.

    The presentation reviewed a range of wave motions with which I am now becoming familiar. Some of these include wave photospheric granular motions inducing horizontal motions or vertical motions inducing p-modes. The work of DeMoortel2012 present observations showing that transverse oscillations are present in a multitude of coronal structures.  It has been observed that inclined fields and mode conversion phenomena enable slow  modes to propagate to the Corona. Pascoe2012 used computational models to demonstrate mode coupling and the spatial damping of propagating kink waves. As well as transverse oscillations, torsional modes have formed an important area of study. Magnetic elements in the G band have allowed the study of the motion of bright points, Jess2006 have suggested that the swirling motions result in the generation of Alfven waves. WedemeyerBohm et al 2012 have taken this further to illustrate how these magnetic tornadoes channel energy into the solar atmosphere. The generation of the propagation of Alfven waves in spicules has been discussed by De Pontieu2011.  Judge2012  consider how typeII Spicules transfer mass and energy to the Corona. Jess2012 compute the chromospheric energy flux of type I spicule oscillations. Evidence for Solar chromospheric fibrils tracing the magnetic field de la Cruz Rodríguez2012. A further example is the demonstration of the magnetoacoustic modes in solar coronal structures resulting from the oscillatory motions of umbral dotsJess2012 achieve this using a combination of numerical modelling and observations.

    Some interesting results were presented illustrating the transfer of energy into the corona by different wave modes including evidence that Chromospheric waves are excited by photospheric wave motions. A rather curious result demonstrated that increasing energy occurs as a result of decreasing period suggesting that high frequency amplitude waves transfer more energy into the Chromosphere. To understand such behaviour it is important to consider the effects of radiative cooling in the Chromosphere, Leenaarts2011 consider just this issue. The 2D Radiative MHD Simulations of Juan Martinez-Sykora2012 illustrate the Importance of Partial Ionization in the Chromosphere.

    In summary
    • Increasing resolution will enable fine scale structure studies.
    • There is an abundance of MHD wave motions at all layers, it is a challenge to link the observational, theoretical and numerical models of these modes.
    • What is happening in the interface region of the solar atmosphere?
    • Insight may be achieved through the synthesis of observations and numerical simulations.
    A number of unanswered questions include:
    • What are the mechanisms for wave generation?
    • What is the total wave flux?
    • What is the mode of transmission of energy from the Chromosphere to the Corona?
    Advances in Solar physics have been achieved using a range of earth based and space borne solar telescopes and spectrometers. A promising development is the Interface Region Imaging Spectrograph (or IRIS) which is a planned space probe to observe the Sun by NASA. It is a NASA Small Explorer program mission to investigate the physical conditions of the solar limb, particularly the chromosphere of the Sun. It will provide information on a large region of the corona from 5-20Mm.

    Saturday, 12 January 2013

    This year has started with a bang!

    After the release of SMAUG I'm helping a couple of users with getting started. We have a lot of work to undertake including documentation and improving usability.

    We continue to test the code and we are currently preparing a model which will run across multiple GPUs. The problem we will consider in this validation exercise is that of modelling the solar atmosphere with the chromosphere acting as a resonant cavity. The solar atmosphere model based on modelling and measurements of VALIIIc and McWhirter, with this representation of the quiet sun our initial tests consider a magnetic field free region so we are using a hydrodynamics model, this is ideal for code testing but the physics is interesting. With a periodic point drivers of periodicity 30 and 300 seconds we can see surface waves propagating along the transition zone. It's very encouraging. The next step will be to extend the height of our model atmosphere and understand how the modes propagate through higher levels in the corona.

    Reports from SDO say that this week an interesting active region has come into view, there is a suggestion that flaring activity may result. The movies from SDO show how the loops evolve and one of the phenomena which is observed is the reconfiguration of field loops.

    Over the last few weeks I've been attending the regular Solar Physics seminars. This week ended with a seminar about failed filament eruptions, these are events which don't result in a full coronal mass ejection.
    Erupting filaments often show a helical structure characteristic of a magnetic flux rope. These helical structures, in addition to being twisted, show evidence of being kinked, the axis of the flux rope exhibits a large-scale writhe. The eruptions may occur as the result of a number of processes

    • kink phenemonena
    • Magnetic Helicty and twisting
    • Reconnection

    During the seminar on failed filament eruptions we were presented with evidence from SDO for twisted flux tubes. Also the characteristic sigmoid pattern after a reconnection event was observed. For the scenario presented in the seminar the enclosing loop had effectively contained the filament eruption.
    Reconnection is a phenomena responsible for a significant quantity of energy release. As suggested by its name it is that process by which magnetic field lines undergo topological reconfiguration. A simple examples is illustrated by the analogue diagram shown above. The cartoon below illustrates a sequence of coronal loops with potential to reconfigure the carpet of magnetic field lines.

    The image below identifies the reconnection events which take place through the life of a coronal loop. Often, during periods of activity these phenomena may be observed in SDO imagery and particularly the AIA171 instrument.