Free Access
Issue
A&A
Volume 546, October 2012
Article Number A23
Number of page(s) 6
Section The Sun
DOI https://doi.org/10.1051/0004-6361/201219268
Published online 01 October 2012

© ESO, 2012

1. Introduction

It is not easy to measure the turbulent magnetic diffusivity ηT at the solar surface. This quantity determines the decay of magnetic structures with scales larger than those of the turbulence. Theoretically, the decay of the magnetic structures should depend on the relation of the magnetic field amplitude to the so-called equipartition value defined by the turbulence. This phenomenon is known as the effect of η-quenching, i.e., the suppression of the eddy diffusivity by the magnetic field.

The simplest realization of η-quenching at the solar surface can be given with two numbers. The decay of active regions after Schrijver & Martin (1990) can be understood with an eddy diffusivity of 1012 cm2/s, while the decay of sunspots with their much stronger fields leads to 1011 cm2/s (Stix 1989). These values are smaller than the value of 3 × 1012 cm2/s, which results from the widely used formula ηT ~ cηurmscorr with the tuning parameter cη ≃ 0.3, the correlation length corr, and parameter values taken close to the surface. Up to now, there was no possibility to measure the turbulent diffusivity on the solar surface for the quiet Sun, where magnetic quenching of this quantity by large-scale magnetic fields is negligible.

Rüdiger et al. (2011) have shown that the combination of a vertical field with driven turbulence in a density-stratified medium leads to an anticorrelation of the cross helicity and the vertical large-scale field, i.e., ⟨u·b⟩ = −ηTBz/Hρ with Hρ as the scale height of the density. If both the cross helicity and the large-scale vertical field are known, then the ratio of the eddy diffusivity and the density scale can be computed. If also the density scale is known from calculated atmosphere models, then fluctuation measurements can be used to calculate the numerical value of the eddy diffusivity for weak fields because if the large-scale magnetic field has only a vertical component and the only vertical gradient is a consequence of the density stratification, then ⟨u·b⟩ ≃ ⟨uzbz⟩. The correlation of the vertical components of flow and field can empirically be obtained by both Doppler measurements and spectropolarimetry.

To estimate the value of the cross helicity, we assume a density scale height of 100 km and write the result in the form (1)where H7 = Hρ/100  km and ηT = 1012  η12 cm2/s. With observations of the left hand side of (1) of about 1 km s-1, one would find ηT of order 1012 cm2/s. This is in line with the results from box simulations of mixed-polarity magnetic fields in the solar surface layer, where values between 1012 cm2/s and 3.4 × 1012 cm2/s were found (Cameron et al. 2011). In the present paper, numerical simulations of stratified magnetoconvection and observational results are discussed, and the theory will be extended to include a vertical stratification of the turbulence intensity. Both the simulations and the observations lead to very similar results for the desired magnetic eddy diffusivity for the quiet Sun, exceeding the value given by (1) by a factor of (only) two.

A simple prediction of this theory is that the ratio (1) does not depend on the sign of the mean magnetic field, i.e., it does not vary from cycle to cycle and (for a dipolar field) from hemisphere to hemisphere. As a consequence, the sign of the cross helicity ⟨u·b⟩ should vary from cycle to cycle and between the hemispheres. Zhao et al. (2011) indeed found indications for a variation from hemisphere to hemisphere in magnetograms and dopplergrams recorded by the Solar and Heliospheric Observatory’s Michelson Doppler Imager in 2000, 2004, and 2007.

2. Mean-field electrodynamics

Let U + u and B + b be the fluctuating velocity and magnetic field with the average values U and B. The scalar correlation between the fluctuations of flow and field, i.e., the cross helicity ⟨u·b⟩, is a pseudoscalar. In the same sense, the cross-correlation tensor ⟨uibj⟩ is a pseudotensor. We are here only interested in its symmetric part (2)As we have shown, the tensor Hij can be finite in presence of a mean magnetic field B and for density-stratified fluids (Rüdiger et al. 2011). Consider these quantities as small enough so that expressions linear in the mean magnetic field influence of these quantities are sufficient. The same may hold for the shear, which influences the (radial) magnetic field components. It is then straightforward to formulate the relation (3)No other formations are possible linear in the mean field B, the stratification vector G, and the shear of the divergence-free mean flow U. For the tensor components, we find (4)and (5)if a box coordinate system (x,y,z) for the latitudinal, azimuthal, and vertical direction is introduced. The z-axis is aligned with the stratification vector, i.e., it represents the radial direction in spherical geometry. The x and y coordinates denote the horizontal directions. Without shear, the correlation Hzz measures the vertical magnetic field and the correlation Hyz measures the azimuthal field. The correlations are also influenced by the shear Uy,z. With the shear included, finite values result for both the correlations (4) and (5), even when the field has only one component. For known values of the correlations, the coefficients and the vertical field of both the azimuthal field and the shear can be computed. We cannot, however, be sure that all the coefficients a....e must be nonzero. First test calculations of Hzz under the presence of horizontal field and shear did not yield finite values of e (Brandenburg, priv. comm.).

The turbulent flow is assumed anelastic, so that div ρu = 0. It is convenient to use the Fourier transformation of the momentum density m = ρu, i.e., (6)and likewise for the fluctuation of the magnetic field.

The spectral tensor of the momentum density that accounts for the stratification of the turbulence to the first-order terms reads (7)where k = (zz′)/2,  κ = z + z′, is the Fourier transform of the local spectrum, (8)so that (9) Derivation of the cross-correlation yields (10)where G = log ρ is the gradient of density and where η is the molecular magnetic diffusivity. Both quantities remain finite in the high-conductivity limit.

From the cross-correlation tensor (10), the cross helicity ⟨u·b⟩ = ηT(G·B) is obtained. From Eq. (10) we find the slightly more complicated expression (13)where G = Gz is the only nonzero radial component of the density-stratification vector. Note the negativity of G. An upwards divergence of the mean field would reduce the effect of density stratification, but for uniform field components, the result is ⟨uzbz⟩ = ⟨u·b⟩.

A real difference between both correlation expressions is, however, due to a possible gradient G′ of the turbulence intensity urms. One easily finds that for vertical fields the turbulence intensity gradient G′ enters expressions for the correlations such as (14)In the bulk of the convection zone, is positive while G is negative. Hence, |⟨uzbz ⟩|>|⟨ u·b⟩| for positive Bz, which is confirmed by the presented simulations (see below).

By elimination of G′ one finds (15)The magnetic eddy diffusivity can thus be determined if the LHS of (15) is calculated from magnetoconvection simulations when the density scale height Hρ is known from numerical models of the solar atmosphere. As only the correlation ⟨uzbz⟩ can directly be observed, one needs a numerical model for the application of the LHS of (15) to derive the eddy diffusivity at the solar surface.

3. Numerical simulations

We perform simulations for a number of different parameter combinations. These parameters include the strength of the imposed vertical field Bz, the viscosity ν, and the magnetic diffusivity coefficient η.

The numerical simulations are done using the Nirvana code, which uses a conservative finite difference scheme (Ziegler 2004). We use Cartesian coordinates. The code solves the equation of motion, (16)the induction equation, (17)the equation of mass conservation, (18)and the equation of energy conservation, (19) In Eqs. (16) and (19), fe is the (external) gravity force and (20)the viscous stress tensor. The total energy density is the sum of the thermal, kinetic, and magnetic energy density: (21) We assume an ideal gas with a constant mean molecular weight μ = 1. The thermal energy density is then (22)with γ = cp/cv = 5/3.

The gas is heated from below and kept at a fixed temperature at the top of the simulation box. Periodic boundary conditions apply at the horizontal boundaries. A homogeneous vertical magnetic field is applied. The upper and lower boundaries are impenetrable and stress-free.

The simulation volume is a rectangular box. The stratification is along the z-coordinate, and it is piecewise polytrophic, with the polytrophic index chosen such that the hydrostatic equilibrium state is convectively stable in the lower and unstable in the upper half of the simulation box. In the following, p denotes gas pressure, ρ mass density, T temperature, g gravity, κ thermal conductivity, and cp the specific heat capacity at constant pressure.

The gas is initially in hydrostatic equilibrium, i.e., (23)where g = const., and the heat flux through the box is vertical and constant, (24)The equation of state is that for an ideal gas. The heat conductivity is constant in the upper and lower layer, respectively, but its values differ between the two layers.

In the dimensionless units, the size of the simulation box is 8 × 8 × 2 in the x, y, and z directions, respectively. The numerical resolution is 512 × 512 × 128 grid points. The stratification of density, pressure, and temperature is piecewise polytrophic, as described in Ziegler (2002). Similar setups have been used by Cattaneo et al. (1991), Brummell et al. (1996), Brandenburg et al. (1996), Chan (2001) and Ossendrijver et al. (2001). The initial state is in hydrostatic equilibrium but convectively unstable in the upper half of the box. The z coordinate is negative in our setup, with z = 0 at the upper boundary. The stable layer thus extends from z = −2 to z = −1, the unstable layer from z = −1 to z = 0. The density varies by a factor 5 over the depth of the box, i.e., the density scale height is 1.2.

thumbnail Fig. 1

Fluctuations of density and temperature in the upper part of the unstable layer at z = −0.05 for Ra = 107 and B0 = 1.

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Figure 1 shows snapshots of the fluctuations of density and temperature for Ra = 107 in a horizontal plane close to the upper boundary. The density is increased at the boundaries of the convection cells and decreased at the center. The opposite is true for the temperature, which is highest at the center of a convection cell and lowest at the boundaries. Vertical velocity is positive, i.e. upwards, at the center and negative, i.e. downwards, at the boundaries. The magnetic field is strongly concentrated in a few small patches which coincide with cell corners, where the gas horizontal flow converges and the vertical flow is downwards.

The initial magnetic field is vertical and homogeneous. We run the simulations until a quasistationary state evolves. Our control parameters are the heat conduction coefficient κ and the Prandtl number Pr = ν/κ. Convection sets in if the Rayleigh number (25)with the density ρ, the specific heat capacity cP, the gravity force g, and the length scale d, exceeds a critical value. The length scale is defined by the depth of the convectively unstable layer, i.e., d = 1. After (1), the correlations and the mean magnetic field always have opposite signs. This has also been confirmed numerically. For positive values of the mean magnetic field Bz, the cross helicity is negative in the unstably stratified layer. If the field polarity is reversed and everything else is left unchanged the cross-correlation becomes positive with the same amplitude.

The velocity field, which is measured in units of cac/100, shows the asymmetry between upwards and downwards motion as characteristic of convection in stratified media. The downwards motion is concentrated at the boundaries of the convection cells and particularly at the corners. The upwards motion fills the interior of the convection cells (see Fig. 2). As it covers a much larger area, the gas motion is considerably slower than in the concentrated downdrafts. The magnetic field shows a similar pattern. The vertical field is concentrated in the areas with downwards motion and weak in the areas with upwards motion. This is the result of field advection, as the total vertical magnetic flux is conserved.

thumbnail Fig. 2

The same as in Fig. 1 but for the fluctuations of the vertical flow and the vertical field.

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thumbnail Fig. 3

Numerical values for the cross helicity ⟨u·b⟩ (left) and the coefficient ⟨uzbz⟩ (right) for weak magnetic field Bz = 10-3. The blue lines denote individual snapshots, and the red lines average over the snapshots shown.

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thumbnail Fig. 4

The same as in Fig. 3 but for B0 = 1.

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Figures 3 and 4 hold for Ra = 107 and for weak and strong magnetic fields. The value of both the Prandtl number and the magnetic Prandtl number is 0.1. The left diagram shows the horizontal average of the cross helicity as a function of the depth, and the right diagram shows the same for the correlation of the vertical components, ⟨uzbz⟩. There is a difference between the two quantities, with the vertical component actually being twice the cross helicity. Equation (15) can thus be written as (26)with (27)The correlations do not vanish abruptly at the bottom of the unstable layer because of overshoot, which affects the upper half of the stable layer. The correlations there are positive and much smaller than in the unstable layer.

The results in Fig. 3 are given in arbitrary units defined by the code. Velocities are given in units of cac/100 with the isothermal speed of sound cac. With an approximate value of cac ≃ 6.6 km s-1 at the optical depth τ = 1 of the Sun, the simulations lead to the cross-correlation velocity ⟨uzbz⟩/Bz ≃ −9 in units of 0.066 km s-1 (Fig. 3, right), i.e., after (26) (28)This value depends only slightly on the magnetic field amplitude for weak fields. For the much stronger magnetic field, Fig. 4 (right) yields the slightly smaller value of 0.81 km s-1.

A characteristic velocity results as the cross-correlation velocity (29)Using the maximal values in Fig. 3 (left), we find Uc ≃ 6 in units of cac/100. Hence, the simulations lead to the cross-correlation velocity Uc ≃ 0.4 km s-1. For the Bz = 1 case (Fig. 4, left), we find Uc ≃ 3 in units of cac/100 or 0.2 km s-1, respectively.

It also makes sense to normalize the cross-correlation in the form (30)which is the ratio of the cross-correlation velocity (29) and the rms velocity of the turbulence. Its numerical value does not depend on the internal units of the code, so that cη is a general and basic result of the simulations. Close to the surface, the maximal numerical value is cη ≃ 0.6 for Bz = 10-3 and cη ≃ 0.3 for Bz = 1. Test calculations for various magnetic fields over many orders of magnitudes show this value as almost uninfluenced by the magnetic-field suppression. Resulting from the overshoot phenomenon at the bottom of the unstable layer, small negative values always appear there. The correlation coefficient (31)for the cross helicity is much smaller than (30) as always (32)similar to the result of Ossendrijver et al. (2001). The relation (31) proves to be true for all amplitudes of the mean magnetic field between 10-5 and 0.1. One finds for all calculations a characteristic correlation coefficient c ≃ 0.1. The Bz = 1 case shows the beginning of the suppression of the fluctuations by the mean magnetic field, which occurs at large values of Bz, resulting in a smaller value of 25 for .

thumbnail Fig. 5

The same as in the right part of Fig. 4 but with reduced numerical resolution of 128 × 128 × 128 (left) and 256 × 256 × 128 (right).

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thumbnail Fig. 6

Numerical values for urms (left) and brms (right) for weak magnetic field, Bz = 10-3 (top) and strong magnetic field, Bz = 1 (bottom). Only the magnetic fluctuations depend on the background field amplitude.

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During the simulations, there are significant temporal fluctuations. The convective instability initially grows exponentially until its saturation, when the system settles in a statistically steady state, but the cross helicity still shows some variations. We therefore average over a certain number of snapshots, typically ten. To test how much the results depend on the numerical resolution, we rerun the Bz = 1 case at the lower resolutions of 256 × 256 × 128 and 1283. Figure 5 shows ⟨uzbz⟩ from those runs. There is a weak dependence on resolution, with higher resolution leading to larger values.

Figure 6 contains all the information about the kinetic and magnetic energies of the magnetoconvection. The rms value of the velocity is hardly influenced by the large-scale magnetic field. In physical units, we find an averaged value of urms ≃ 0.1cac ≃ 0.66 km s-1. In contrast, the magnetic energy strongly depends on the applied magnetic field. In dimensionless units, it is brms/urms ≃ 0.6Bz in both cases, which leads to (33)in physical units. At the top of the convection zone, we find very small contributions of the magnetic energy for Bz = 1 Gauss, while for 1000 Gauss there is almost equipartition.

Table 1

Results from the analysis of the SST and Hinode data.

4. Observations

It is difficult to empirically determine the cross helicity ⟨u·b⟩ at the solar surface because it is hard to retrieve the horizontal flows and magnetic field components from observations. We have, however, the possibility to use the relation ⟨u·b⟩ ≈ 0.5⟨uzbz⟩, known from the above numerical simulations. The vertical flow speed and magnetic field component can be determined with much better accuracy. Then using Eq. (26), we can determine the cross-helicity velocity from the observations.

For this purpose, we have analysed two datasets containing observations of the quiet Sun at disk center, where the line-of-sight coincides with the local vertical. Data from the Crisp Imaging Solar Polarimeter (CRISP) instrument on the Swedish 1-m Solar Telescope (SST) cover the 6302.5 Å Fe i spectral line with 12 equidistant wavelength positions at 48 mÅ steps and a continuum point. They have a pixel scale of and a total field-of-view of about 60′′  ×  60′′. The second dataset is from the spectropolarimeter on the Solar Optical Telescope (SOT) of Hinode. It covers both the 6301.5 and 6302.5 Å Fe i lines, has a pixel scale of , and a total (scanned) field of view of 164′′  ×  328′′.

The line-of-sight velocity and magnetic field data for the Hinode observations were taken from the level two data products available online1. Magnetic field strengths have been converted to fluxes by taking the filling factor into account. The SST data were inverted using the lilia inversion code (Socas-Navarro 2001). Velocities were calibrated using the convective blueshift determined by de La Cruz Rodríguez et al. (2011). More details on these two datasets can be found in Schnerr & Spruit (2011).

We show the results for these datasets in Table 1. The cross-helicity velocity (uη) as determined from the SST data is somewhat higher than that from the Hinode data. This is at least partly due to the lower resolution of Hinode as compared to the SST. If we rebin the SST data to a lower resolution, the cross-helicity velocity decreases (see Table 1) because the strongest fields and flows are smoothed out.

The coefficient from the Hinode and SST data is 521.3 and 163.5 respectively, which is larger than the value of 50 found in the simulations. This indicates that the effective magnetic Reynolds number in the simulations is smaller than in the solar convection zone.

5. Conclusions

We have shown that nonrotating turbulence at the top of the solar convection zone under the influence of a vertical magnetic field forms a finite cross helicity. The only condition is the existence of a vertical stratification of density and/or turbulence intensity. While the effect would not appear within the Boussinesq approximation, it exists in the high-conductivity limit, i.e., for sufficiently large magnetic Reynolds numbers.

In our understanding, the cross helicity is anticorrelated to the mean radial magnetic field, i.e., (34)For an oscillating dipolar background field, the sign of the cross helicity differs for both hemispheres and also from cycle to cycle.

The theory can also be used to measure the magnetic diffusivity if the cross helicity is known by observations. In order to find the cross helicity, one has only to correlate observed flow fluctuations with observed magnetic fluctuations.

The anticorrelation (34) for density-stratified turbulence has been established by Rüdiger et al. (2011) for a model of numerically-driven turbulence. In the present paper, buoyancy-driven magnetoconvection has been simulated in a box with the Nirvana code. We find that such a turbulence also fulfills the relation (34). The correlation coefficient (30) takes the value of 0.6 for the weak magnetic field Bz = 10-3 and 0.3 for the stronger field Bz = 1. The ratio (32) of the magnetic fluctuations to the applied magnetic field is always of the order five.

We have also shown that for density-stratified turbulence the identity ⟨u·b⟩ = ⟨uzbz⟩ holds. So far, solar observations can only measure the correlation ⟨uzbz⟩. The numerical simulations,

however, always lead to the result ⟨uzbz⟩ ≃ 2⟨u·b⟩, so that the observed value of ⟨uzbz⟩ would overestimate the actual cross helicity by a factor of two. The reason is the vertical stratification of the turbulence intensity, which at the top of the convection zone is antiparallel to the density stratification. Hence, both the correlations ⟨uzbz⟩ and ⟨u·b⟩ are reduced but not by the same amount.

With ⟨uzbz⟩ ≃ 2⟨u·b⟩, the value of uη can be computed by using Eq. (26). The numerical simulations lead to uη ≃ 1 km s-1 and uη ≃ 0.8 km s-1 respectively for the two cases studied. This result is well confirmed by the observations that lead to values between 0.7 km s-1 (Hinode) and 1.3 km s-1 (SST). To estimate the value of the eddy diffusivity at the solar surface, we shall assume a density scale height of 100 km and find values close to ηT ≃ 1012 cm2/s for the eddy diffusivity at the surface of the quiet Sun.


Acknowledgments

We gratefully acknowledge Axel Brandenburg (Stockholm) for motivating discussions and numerical support.

References

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All Tables

Table 1

Results from the analysis of the SST and Hinode data.

All Figures

thumbnail Fig. 1

Fluctuations of density and temperature in the upper part of the unstable layer at z = −0.05 for Ra = 107 and B0 = 1.

Open with DEXTER
In the text
thumbnail Fig. 2

The same as in Fig. 1 but for the fluctuations of the vertical flow and the vertical field.

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In the text
thumbnail Fig. 3

Numerical values for the cross helicity ⟨u·b⟩ (left) and the coefficient ⟨uzbz⟩ (right) for weak magnetic field Bz = 10-3. The blue lines denote individual snapshots, and the red lines average over the snapshots shown.

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In the text
thumbnail Fig. 4

The same as in Fig. 3 but for B0 = 1.

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In the text
thumbnail Fig. 5

The same as in the right part of Fig. 4 but with reduced numerical resolution of 128 × 128 × 128 (left) and 256 × 256 × 128 (right).

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In the text
thumbnail Fig. 6

Numerical values for urms (left) and brms (right) for weak magnetic field, Bz = 10-3 (top) and strong magnetic field, Bz = 1 (bottom). Only the magnetic fluctuations depend on the background field amplitude.

Open with DEXTER
In the text

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