A&A 466, 1131-1144 (2007)
DOI: 10.1051/0004-6361:20067017
R. Rezaei1 - R. Schlichenmaier1 - C. A. R. Beck1,2 - J. H. M. J. Bruls1 - W. Schmidt1
1 - Kiepenheuer-Institut für Sonnenphysik, Schöneckstr. 6, 79 104 Freiburg, Germany
2 -
Instituto de Astrofísica de Canarias (IAC), 38 205 La Laguna, Espain
Received 22 December 2006 / Accepted 25 January 2007
Abstract
Aims. We investigate the relationship between the photospheric magnetic field and the emission of the mid chromosphere of the Sun.
Methods. We simultaneously observed the Stokes parameters of the photospheric iron line pair at 630.2 nm and the intensity profile of the chromospheric Ca II H line at 396.8 nm in a quiet Sun region at a heliocentric angle of 53.
Various line parameters have been deduced from the Ca II H line profile. The photospheric magnetic field vector has been reconstructed from an inversion of the measured Stokes profiles. After alignment of the Ca and Fe maps, a common mask has been created to define network and inter-network regions. We perform a statistical analysis of network and inter-network properties. The H-index is the integrated emission in a 0.1 nm band around the Ca core. We separate a non-magnetically,
,
and a magnetically,
,
heated component from a non-heated component,
in the H-index.
Results. The average network and inter-network H-indices are equal to 12 and 10 pm, respectively. The emission in the network is correlated with the magnetic flux density, approaching a value of
pm for vanishing flux. The inter-network magnetic field is dominated by weak field strengths with values down to 200 G and has a mean absolute flux density of about 11 Mx cm-2.
Conclusions. We find that a dominant fraction of the calcium emission caused by the heated atmosphere in the magnetic network has non-magnetic origin (
pm,
pm). Considering the effect of straylight, the contribution from an atmosphere with no temperature rise to the H-index (
pm) is about half of the observed H-index in the inter-network. The H-index in the inter-network is not correlated to any property of the photospheric magnetic field, suggesting that magnetic flux concentrations have a negligible role in the chromospheric heating in this region. The height range of the thermal coupling between the photosphere and low/mid chromosphere increases in presence of magnetic field. In addition, we demonstrate that a poor signal-to-noise level in the Stokes profiles leads to a significant over-estimation of the magnetic field strength.
Key words: Sun: photosphere - Sun: chromosphere - Sun: magnetic fields
The chromospheric heating mechanism is one of the main challenges of solar physics (Narain & Ulmschneider 1996, and references therein). The core emission of the Ca II H and K lines is an important source of radiative losses in the chromosphere. Moreover, this emission is an important tool to study the temperature stratification and the magnetic activity of the outer atmosphere of the Sun and other stars (e.g., Schrijver & Zwaan 2000). Most of the observational studies based on these lines use either the calcium intensity profile (e.g., Cram & Damé 1983; Lites et al. 1993) or combinations of calcium filtergrams and magnetograms (e.g., Berger & Title 2001). From the observation of the calcium spectrum alone, it is not possible to distinguish between the magnetic and non-magnetic heating components. Combining calcium filtergrams with magnetograms allows to separate those components, but the spectral information is lost. Simultaneous observations that allow to reconstruct the magnetic field and record the spectrum for the Ca II H line are rare (e.g., Lites et al. 1999) and only available at lower spatial resolution. Therefore, it is not surprising that none of the present theories, mechanical and Joule heating, was confirmed or rejected observationally (Socas-Navarro 2005; Fossum & Carlsson 2005).
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Figure 1: From left to right: a) the Fe I 630 nm continuum intensity, b) the calcium wing intensity at 396.490 nm which was calibrated to FTS data (Stenflo et al. 1984), c) the outer calcium wing intensity (W1), d) the inner calcium wing intensity (W3), e) the H-index, f) the masks which separate network (white) from the inter-network (gray), and g) the magnetic flux density obtained from the inversion. We did not use the black region between the network and inter-network. Each small tickmark is 1 arcsec. Note that sampling in x and y directions are different. |
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The POlarimetric LIttrow Spectrograph (POLIS, Schmidt et al. 2003; Beck et al. 2005a) was designed to provide co-temporal and co-spatial measurements of the magnetic field in the photosphere and the Ca II H intensity profile. We use POLIS to address the question of the chromospheric heating mechanism, by comparing properties of network and inter-network in photosphere and chromosphere by a statistical analysis. With the information on the photospheric fields, we separate the contributions of the magnetically and non-magnetically heated component. For the first time we study the correlation of the chromospheric emission with the corresponding amplitude/area asymmetry and Stokes-V zero-crossing velocity at the corresponding photospheric position. We also investigate the magnetic field strength distribution of the inter-network to check whether it consists of weak fields (e.g., Faurobert et al. 2001; Collados 2001) or kilo-Gauss fields (e.g., Sánchez Almeida et al. 2003a,b).
Observations and data reduction are explained in Sects. 2 and 3. Histograms of the obtained parameters are presented in Sect. 4. Correlations between the chromospheric and photospheric parameters are addressed in Sect. 5. The heating contributions are elaborated in Sect. 6. Discussion and conclusions are presented in Sects. 7 and 8, respectively. Details of the magnetic field parameters in the inter-network are discussed in Appendix A. The straylight contamination and calibration uncertainties for the POLIS Ca channel are estimated in Appendix B. An overview of some of our results also appears in Rezaei et al. (2007).
All maps were recorded with a slit width of 0.5 arcsec, a slit height of 47.5 arcsec, and an exposure time per slit position of 4.92 s. The scan extension for the first three maps was 40.5, 55.5, and 20.5 arcsec, while it was 25.5 arcsec for the remaining ten maps.
The Ca II H line and the visible neutral iron lines at 630.15 nm, 630.25 nm, and Ti I 630.38 nm were observed with the blue (396.8 nm) and red (630 nm) channels of POLIS. The spatial sampling along the slit (y-axis in Fig. 1) was 0.29 arcsec. The scanning step was 0.5 arcsec for all maps. The spectral sampling of 1.92 pm for the blue channel and 1.49 pm for the red channel leads to a velocity dispersion of 1.45 and 0.7 km s-1 per pixel, respectively. The spectropolarimetric data of the red channel were corrected for instrumental effects and telescope polarization with the procedures described by Beck et al. (2005b,a).
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Figure 2: Sample averaged calcium profile of one of the maps (the average profile is similar for all other maps). The bands are explained in Table 1. |
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Table 1: The definition of the characteristic parameters of the Ca II H profile for the peak sample (upper part) and the band sample (lower part) (see also Fig. 2). Wavelengths are in nm.
Figure 1 displays an overview of one of the thirteen maps after spatial alignment. The map (a) shows the Fe I 630 nm continuum intensity normalized to the average quiet Sun intensity. The map (b) shows the Ca II H wing intensity, taken close to 396.490 nm (for a definition of line parameters see Fig. 2 and Table 1). These two maps were used for the spatial alignment of the red and blue channels. The next two maps (c and d) show the intensities in the outer and inner wings (W1, respectively, W3). The inner wing samples a wavelength band close to the core; hence, it is more influenced by the line-core emission and shows higher contrast of the network than the outer wing. The next map (e) is the H-index, i.e., the intensity of the calcium core integrated over 0.1 nm (cf. Table 1). The network features appear broadest and show the highest contrast in this map. The map (f) demonstrates network and inter-network masks (Sect. 3.1). The map (g) shows the magnetic flux density obtained from the spectro-polarimetric data (Sect. 3.2).
For each map, we created a mask to distinguish between the network and inter-network regions (Fig. 1f). This was done manually on the basis of the magnetic flux and the H-index. We did not use the black region in Fig. 1f which separates network from inter-network. To study structures in these two regions, we define two statistical samples:
Table 1 lists the characteristic parameters we define for each calcium profile:
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Figure 3:
Scatter plot of the
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Figure 4:
Histograms of the intensity parameters for the network (thick) and
inter-network (thin), using the band sample: a) the H-index, b)
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An inversion was performed using the SIR code (Ruiz Cobo & del Toro Iniesta 1992). We used the same setup as in Bellot Rubio & Beck (2005) and Beck et al. (2006): a two-component solar atmosphere model with one magnetic and one field-free component. Additionally, a variable amount of straylight was allowed for. This accounts for unresolved magnetic fields inside each pixel. We did not consider any gradient for the atmospheric parameters except for the temperature. The inversion yields the magnetic field vector and an estimate for the magnetic filling fraction. The flux density map (magnetic flux per pixel, Fig. 1g) is based on the inversion results.
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Figure 5: Histograms of the magnetic field parameters for the network (thick) and inter-network (thin), using the peak sample: a) the absolute magnetic flux density, b) the field strength, c) Stokes-V amplitude, d) Fe I 630.25 nm V velocity, e) the amplitude asymmetry, and f) the area asymmetry. |
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Figure 6:
Upper panels: correlation between the
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The V velocity distribution is shown in Fig. 5d. Although it has a peak around zero both for the network and inter-network, the inter-network shows a larger fraction of high-velocity V profiles. The amplitude and area asymmetries in the network peak at small positive values of 10% and 3%, respectively (Figs. 5e and f). In contrast, there is no tendency to positive or negative values for the inter-network asymmetries. This is in agreement with other studies of quiet Sun magnetic fields (Sigwarth et al. 1999; Khomenko et al. 2003; Sigwarth 2001).
The (V/R) ratio is an indicator for the H3 line-core position (Rutten & Uitenbroek 1991; Cram & Damé 1983).
There is a strong correlation between the (V/R) ratio
and (H3): the redshifted
and blueshifted calcium profiles correspond to (V/R) ratios
smaller and larger than one, respectively (Fig. 6, lower right panel)
.
We find a similar correlation between the emission strength and the (V/R) ratio:
the larger the (V/R) ratio, the larger the emission strength.
The lower left panel of Fig. 6 shows the correlation between the outer wing intensity and the H-index.
While there is a correlation in the network (gray), there is no significant correlation in the inter-network (Table 2).
We calculated the correlation coefficient between the H-index and all
intensities in the line wing to investigate this difference in more
detail (Table 2). In the network, there is a stronger correlation than
in the inter-network for all the wing bands considered.
We return to this point in Sect. 7.
There is a strong correlation between the amplitude and area asymmetries, both in the network and inter-network (Fig. 7, top left panel). In contrast, there is no correlation between either amplitude or area asymmetries and the V velocity (bottom left panel, Fig. 7). These are typical properties of the quiet Sun magnetic field (e.g., Sigwarth et al. 1999). The right panels of Fig. 7 show scatter plots of the V asymmetries versus V amplitude. There is a tendency for high amplitude V signals to have small asymmetries. On the other hand, histograms of the asymmetries peak at a positive value. Implications of these two findings will be discussed in Sect. 5.3.2.
Table 2: Correlation coefficients between the band intensities and the H-index in the network and inter-network using the band sample.
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Figure 7: Top left: scatter plot of the amplitude vs. area asymmetries. Bottom left: scatter plot of the area asymmetry vs. the V velocity. Right panels: scatter plots of the amplitude and area asymmetries vs. the Fe I 630.25 nm V amplitude. Gray and black show the network and inter-network, respectively. |
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The upper panel of Fig. 8 shows the relation between the photospheric magnetic flux and the
chromospheric emission for the network.
For lower flux values (<100 Mx cm-2), there is a clear increase of emission with flux.
However, for higher
magnetic flux densities, the H-index increases slowly.
To reproduce the observed relation, we utilize a power law fit to the data,
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(1) |
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Figure 8:
Upper panel: correlation between the H-index and the absolute
magnetic flux density. Gray is the original data and black is the binned data: each point is average of 25 points.
The middle curve shows a fit of a power law to the original data (Eq. (1)), the other two curves give the
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Figure 9: Correlation between the H-index and amplitude/area asymmetry and V velocity. Plusses and squares show network and inter-network, respectively. The binning method is similar to Fig. 8. |
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The calcium core emission in the inter-network does not correlate with the magnetic flux, at least with the concentrated magnetic field within the Zeeman sensitivity (Fig. 8, lower panel). The value of the offset, c, of the fit to the network (Table 3) for low magnetic flux densities is around 10 pm, which is consistent with the average H-index value of all the inter-network profiles (10 pm for the peak sample). The value of the c parameter changes for different thresholds of the flux density (Table 3). Implications of this finding for the basal flux are discussed in Sect. 7.
Table 3: The parameters of the fit to Eq. (1) to the network data. The first column is the threshold for the magnetic flux density (Mx cm-2). By increasing the threshold, we avoid the inter-network intrusions. For a comparison to the peak inter-network flux density, see Appendix A.
There are different behaviors for the network and inter-network calcium core emission with respect to V velocity and asymmetries (Fig. 9). The core emission peaks at small positive amplitude/area symmetry in the network. This is caused by strong V profiles that show small asymmetries (Fig. 7, right panels). The highest H-index in the network corresponds to almost zero V velocity. In its diagram (Fig. 9), the slope of the left branch (blueshift) is smaller than the right one. Moreover for the strong upflow or downflow in the magnetic atmosphere in the network, the H-index decreases significantly. The H-index in the inter-network does not depend on any parameter of the Stokes-V profile (Fig. 9).
Considering the fact that the amplitude and area asymmetries strongly correlate with each other (top left panel, Fig. 7), a natural consequence is that the left (right) branch of the scatter plot of the H-index vs. amplitude asymmetry corresponds to the left (right) branch of the scatter plot of the H-index vs. area asymmetry (Fig. 9). However, if we compare only negative or positive branches of the scatter plots of the H-index vs. V velocity and asymmetries in the network, we realize that the left branch in the V velocity does not correspond to the similar branch in the asymmetry plots. The left diagrams in Fig. 10 consider only the profiles with a negative V velocity, while in the right panels profiles with a negative area asymmetry are shown. A positive V velocity may thus correspond to a positive or negative area asymmetry. There is no relation between the amplitude/area asymmetries and the V signal in the inter-network (right panels, Fig. 7).
The H-index includes the H1 (the minima outside the emission peaks), H2, and H3 spectral regions. So its formation height extends from the higher photosphere to the middle chromosphere.
The Ca II H & K lines are one of the main sources of the chromospheric radiative loss.
We use the H-index as a proxy for the emission in the low/mid chromosphere, and
assume a linear relation between the H-index and the chromospheric radiative loss.
To derive the contribution of magnetic fields to the chromospheric emission, we
added the H-index of all points with magnetic flux above a given flux threshold.
The fractional H-index, ,
is then defined by normalizing this quantity to the total
H-index of all points in the field of view
:
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Figure 10: In the left column diagrams, only the points with a negative V velocity are plotted. In the right panels, only points with a negative area asymmetry are plotted. Pluses and squares show network and inter-network, respectively. |
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Table 4:
Fractional distribution of the total H-index, ,
for different magnetic flux thresholds.
The average H-index is 12.0 and 9.8 pm in the network and inter-network, respectively.
However, it contains some contributions from the photosphere (outside emission peaks) and a cool
chromosphere (without temperature rise).
Therefore, we decompose the observed calcium profiles in a
heated and a non-heated component,
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(2) |
We use the Holweger & Müller (1974) model atmosphere as a proxy for those parts of the solar atmosphere that are not affected by any heating process. This model is very similar to a theoretical radiative equilibrium model, but it was not constructed to strictly satisfy this condition. We prefer this semi-empirical model over a theoretical radiative equilibrium model because it results in slightly better reproduction of the Ca II H line wings in a standard disk-center atlas. We note, however, that 3D (M)HD simulations of the solar atmosphere commonly produce regions in the upper photosphere and chromosphere with temperatures significantly lower than in the Holweger-Müller model. These structures, however, are generally so small that their evolution is strongly influenced by their surroundings, so that they cannot serve as a model for those parts of the solar atmosphere that are not influenced by any heating processes. That means that they are not suitable for determining the absolute minimum value of the H-index.
We perform a consistent NLTE radiative transfer computation in plane-parallel
geometry, including effects of partial frequency redistribution, in order to
estimate the non-heated component of the H-index.
We added an isothermal hydrostatic extension up to a height of 2000 km
to the model in order to be able to compute the core parts of the Ca II H & K lines.
This may seem an arbitrary choice, but considering that the computed
intensity is essentially zero at the line core for models with such low
temperatures in the upper atmosphere, its influence on the H-index
is very small. This model atmosphere and the obtained profile are
similar to the COOLC model of Ayres et al. (1986) and the corresponding profile.
We find a value of
pm at disk center.
At a heliocentric angle of 53
,
the appropriate value is
pm.
We find an upper limit of 12% for the straylight contamination of the
observed profiles (cf. Appendix B).
Therefore, we decompose the average H-index in the inter-network (
pm) to obtain
the non-magnetically heated component,
:
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(3) |
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(4) |
The correlation between the inner wing intensity and H-index is strong, both in the network and inter-network (Table 2). However, the correlation between the outer wing intensity and the H-index is significant only in the network (Fig. 6, lower left panel). This indicates that the height range of the thermal coupling between the photosphere and low/mid chromosphere increases in presence of a magnetic field. There are also suggestions that this coupling extends to the upper chromosphere (Rauscher & Marcy 2006).
The rate of increase of the chromospheric emission vs. magnetic flux reduces for strong magnetic flux densities. This implies that for large flux concentrations, either the filling factor or the magnetic field strength saturates. There are indications that it is the filling factor: the available space is completely filled by expanding flux tubes (Hammer 1987; Schrijver et al. 1996; Saar 1996). In contrast, a lower limit of chromospheric emission also exists. In the inter-network, no relation between emission and the photospheric fields is found (Figs. 8 and 9). The average value of the H-index in the inter-network of around 10 pm corresponds to the offset, c, Eq. (1). This reflects a constant contribution to the H-index which is present even without photospheric magnetic flux, in agreement with Schrijver (1987,1995) who argued that the basal flux does not depend on the magnetic activity. The basal flux contains two components: the non-heated (cf. Sect. 6) and the non-magnetically heated contributions. The non-heated contribution depends on the temperature stratification, so that we speculate that the non-magnetically heated component has an inverse dependence on the temperature stratification.
Figure 8 (upper panel) shows that there is a variable lower limit for the H-index versus the
magnetic flux density.
In contrast, the upper limit is less clearly defined, and would be in
agreement with a constant maximum value independent of the amount of magnetic flux.
There are also similar behaviors for the upper and lower limits in the scatter plots of
the H3,
,
and
versus the magnetic flux density.
This is similar to upper and lower limits of the basal flux of stars vs. the color (B-V) where
the lower boundary changes but the upper one is almost constant (Fawzy et al. 2002b, their Fig. 2).
The situation for the inter-network is different: both the upper and lower limits are
independent of the magnetic flux density.
The fact that the non-magnetic chromospheric heating contributes significantly to
the chromospheric energy balance is also found in the relative contributions of mainly
field-free inter-network and network areas to the total emission in the field of view.
Magnetic flux densities above 50 Mx cm-2 add
only 20% to the total emission, which is also seen in the ratio of the mean H-index
in the network and inter-network, 12.0 /
(Table 4).
There may be some mixed polarity fields below our polarimetric detection limit which may influence the ratio of the magnetically heated to the non-magnetically heated component.
The magnetically heated component is related to Stokes-V profiles with non-zero area and amplitude asymmetries (Figs. 5 and 9). A possible explanation for this finding would be the absorption of upward propagating acoustic waves, generated by the turbulent convection, by the inclined fields of expanding flux tubes. This would imply that the energy is deposited at the outer boundary or in the canopy of flux concentrations rather than in the central, more vertical, part. It is mainly because the non-magnetic cutoff frequency is lowered at the boundary of the flux tubes, where the field lines are inclined. This was first predicted by Suematsu (1990) and recently achieved some observational support (Jefferies et al. 2006; Hansteen et al. 2006). Figure 9 indicates that the maximum observed H-index has non-zero V asymmetry. These asymmetric V profiles are consistent with the case when the line of sight passes through the canopy of a magnetic element or through a flux tube axis (positive and negative asymmetries respectively, Steiner 1999). Therefore our finding supports Suematsu (1990).
There is no correlation between the H-index and magnetic field parameters in the inter-network.
Therefore, we conclude that the magnetic field has a negligible role in
the chromospheric heating in the inter-network.
On the other hand, the H-index is a power law function of the magnetic flux density
in the network with a power index of 0.3.
The average H-index observed in the network and inter-network are 10 and 12 pm, respectively.
We find a non-magnetic component in the network (based on Eq. (1)) of about 10 pm which
shows the consistency of our analysis: the non-magnetic part of the
network H-index is equal to the H-index of the inter-network.
A NLTE radiative transfer calculation, using the Holweger-Müller model atmosphere, indicates that the non-heated component of the H-index, emerging from a cool chromosphere, is about 5.9 pm. Comparison of this non-heated component and the average H-index in the inter-network has two implications: a) some of the observed chromospheric emission does not originate from a hot chromosphere, and b) the non-magnetically heated component is about 50% larger than the magnetically heated component. From this, we conclude that the non-magnetically heated component has a larger contribution in the chromospheric radiative loss than the magnetically heated component, both in the network and inter-network.
In our statistical ensemble,
spatial positions with strong magnetic field ( 50 Mx cm-2) contribute about 20%
of the total H-index. Correlations and histograms of the different intensity bands in the Ca II H spectrum
indicate that above a magnetic threshold, photosphere and low/mid chromosphere are thermally
coupled. Moreover, our findings are consistent with the idea that the energy transfer in a flux tube
has a skin effect: the energy transfer is more efficient in the flux tube boundary (canopy)
than at its center (axis).
For the first time, we find a magnetic field distribution in the inter-network using visible lines
which is similar to results inferred from infrared lines with larger Zeeman splitting.
The distribution function increases with decreasing field strength. The peak at 200 G is due to the detection
limit of the polarization signal. The average
and distribution peak of the absolute flux density (of magnetic profiles)
in the inter-network are 11 and 4 Mx cm-2, respectively.
We conclude that the combination of high spatial resolution and polarimetric accuracy is sufficient
to reconcile the different results on field strength found from infrared and visible lines.
Acknowledgements
The principal investigator (PI) of the ITP observing campaign was P. Sütterlin, Utrecht, The Netherlands. We wish to thank Reiner Hammer, Oskar Steiner, Thomas Kentischer, Wolfgang Rammacher, and Hector Socas-Navarro for useful discussions.The POLIS instrument has been a joint development of the High Altitude Observatory (Boulder, USA) and the Kiepenheuer-Institut. Part of this work was supported by the Deutsche Forschungsgemeinschaft (SCHM 1168/8-1).
We achieved a magnetic field distribution for the inter-network using the Fe I 630 nm pair that is similar to the infrared studies (Fig. 5b). We elaborate on the effects of spatial resolution and signal-to-noise ratio on this finding. We used the Kiepenheuer Adaptive Optics System (KAOS) to improve spatial resolution and image stability (von der Lühe et al. 2003). It provides stable sharp images in a certain field of view, and therefore allows us to increase the exposure time to have better signal-to-noise ratio.
Figure A.1 is a close-up view of the inter-network magnetic flux distribution.
The lowest flux detected is 0.05 Mx cm-2. The distribution of the
flux decreases sharply for fluxes below
Mx cm-2, which is due
to the detection limit of the polarization signal. The small fluxes still detected
give us, however, confidence that the magnetic field strength distribution of Fig. 5b
is reliable as a consequence of the combination of high spatial resolution and good signal-to-noise ratio.
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Figure A.1:
Close up view of the flux density distribution for the inter-network.
At a value of ![]() |
To investigate the influence of noise on the inversion results,
we compare inversion results of the original datasets with and without adding noise.
In order to create "noisy'' data sets, the rms value of the noise in the original data
was increased to
,
two times larger than the original value.
The spatial sampling was identical to the original data.
We emphasize that inversion was done using the same assumptions and initial model atmosphere.
Figure A.2 compares two sets of full Stokes profiles (black)
and the inversion fit (red). The upper four panels are the original profiles and the lower four
panels are the same profiles with noise. The fits and retrieved model atmospheres differ significantly.
Especially important is the fact that the noisy inversion leads to a magnetic field of 1.5 kG.
To compare the field strength obtained from the original and noisy datasets, their
histograms are shown in Fig. A.3. The distribution of the original field strength
peaks around 200 G (black histogram), while for the noisy data, it has a clear shift
toward higher values with a peak at 0.8 kG (red histogram).
This emphasizes the role of the noise in the existing discrepancy between visible and
infrared measurements (Bellot Rubio & Collados 2003).
In brief, we find that higher spatial resolution of these observations along with low noise data
is an important step toward resolving disagreements between the visible and infrared polarimetric measurements.
The visible Fe I 630 nm pair shows higher Stokes amplitudes for small magnetic fields than the
infrared lines of Fe I 1.56 m. Hence, with an equal amount of noise in observational data,
the visible lines provide a higher signal-to-noise ratio than infrared lines.
Therefore, in contrast to Martínez González et al. (2006), we find the visible Fe I 630 nm pair to be a proper
tool to investigate the inter-network magnetic field. We ascribe the present disagreements
to the low spatial resolution of previous observations and different signal-to-noise ratios.
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Figure A.3: Histograms of the inter-network field strengths of one of the maps based on the inversion of original data (black) and noisy data (red). |
The POlarimetric LIttrow Spectrograph (POLIS) was designed to facilitate co-temporal and co-spatial measurements of the vector magnetic field in the photosphere and the Ca II H intensity profile (Schmidt et al. 2003; Beck et al. 2005a). We try to elaborate on the uncertainties in the intensity calibration, because it is critical for the estimates of the various contributions (straylight, non-magnetic heating, etc.) to the observed calcium profiles. Observations are discussed in Sect. B.1. We study the linearity of the CCD camera for different light levels in Sect. B.2. Then, we investigate instrumental and solar straylight contributions in Sect. B.3. In Sect. B.4, we perform standard error propagation to quantify uncertainties in our final calibrated spectra.
All spectra and profiles in this Appendix are dark subtracted raw data. No other calibration process like flat fielding was applied. In order to show intrinsic noise in the data, we did not spatially average darks or spectra.
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Figure B.3: Left: a single Ca II H spectrum of the umbra and quiet sun profiles in a sunspot. Asterisks show spectral bands that were used to estimate the true signal (see Sect. B.3.2). Right: close up view of the quiet sun and umbral profiles near the calcium line core. The true signal (about 5% in the wing close to the Ca core according to Wallace et al. (2000) was subtracted from the umbral profile. The red curve shows 12% of the quiet sun profile. The two cyan curves show an uncertainty interval of one sigma (1%). |
Figure B.2 shows a slit spectrum of the umbra. To keep the sample profiles as uniform as possible, we chose two identical positions along the slit (at about 30 arcsec, Fig. B.2). The spectrum of the quiet sun and the minimum spectrum of the umbra are shown in Fig. B.3 (left panel).
Wallace et al. (2000) presented a Fourier transform spectral atlas of the umbra.
In this atlas, the wing close to the umbral core
has an intensity of 5% in units of the umbral continuum,
while the wing close to the Fe I, Ti I, and Ni I lines at about
396.5 nm has an intensity of 38%. In our umbral profile,
this continuum has an intensity of about 80 counts (asterisks, Fig. B.3, left panel),
whereas close to the core the profile has about 45 counts.
Assuming no straylight in the atlas profile, only 11 counts (22% of the observed valued)
should be measured in the core. Thus, 78% of the observed signal have to be due
to straylight.
To reproduce 78% of the umbral profile, a straylight contribution of around 12% is needed.
If we assume that there is no real calcium signal in the umbra (which is wrong) and
take the observed calcium profile of the umbra as pure scattered light,
the ratio changes from 12% to 15%.
Therefore, uncertainties about the scattered light in the atlas profile have minor importance.
A close-up view of the spectral region close to the calcium core is shown in Fig. B.3 (right panel).
For the reason mentioned above,
in the right panel of Fig. B.3, the umbral profile was labeled umbra (78%).
The red curve shows 12% of the quiet sun profile. The two cyan curves also show
an uncertainty interval of one sigma (1%). Figure B.4 shows the
ratio of the umbra to the quiet sun profile. If we calculate mean and rms of the whole
spectrum, it gives a value of 10 3% while considering only on the wing band
close to the calcium core, we obtain 12
1%. We take this value as an upper limit
for the total straylight in the POLIS calcium channel.
In Sect. B.3.1, we concluded that the instrumental straylight is about 1%.
Hence, most of the obtained straylight originates from
the scattered solar light in the telescope and the earth atmosphere.
There are three noise sources in a CCD camera (Mullikin et al. 1994):
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Figure B.7: A sample calcium profile. We use the calcium core minimum counts to estimate the photon noise. In this case, the signal-to-noise ratio at the calcium core is more than 10. |
To investigate the thermal noise in the dark current, one has to
remove small systematic offsets between CCD columns
(probably due to thermal fluctuations of the AD system) before
considering statistical measures, e.g., rms, of a dark profile (a CCD row).
As it is shown in Fig. B.6, for a dark frame with 6 accumulations,
the resulting rms (after removing systematics)
is far less than the standard deviation of the whole map including a large-scale variation.
It provides an estimate for the thermal noise in the camera.
So the rms of the dark frame is about
counts.
The total noise in the recorded data,
,
consists of three components:
![]() |
(A.1) |
The relative intensity of the core to wing is defined as:
![]() |
(A.2) |
![]() |
(A.3) | ||
![]() |
(A.4) |
![]() |
(A.5) |