Recent advances in measuring chromospheric magnetic fields in the He i 10830 Å line

Advances in Space Research 39 (2007) 1734–1740 www.elsevier.com/locate/asr

Recent advances in measuring chromospheric magnetic fields ˚ line in the He I 10830 A A. Lagg

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Max-Planck-Institut fu¨r Sonnensystemforschung, Max-Planck-Str. 2, 37191 Katlenburg-Lindau, Germany Received 31 December 2006; received in revised form 22 March 2007; accepted 30 March 2007

Abstract During the last decade advances in instrumentation, atomic physics and modeling have greatly improved the access to the chromospheric magnetic field vector. High sensitivity polarimeters like the Tenerife Infrared Polarimeter (TIP2, VTT) or the Spectro-Polarimeter ˚ triplet. The simultafor Infrared and Optical Regions (SPINOR, HAO) lead to reliable Zeeman measurements using the He I 10830 A ˚ line provides additional information on the structure of the underlying photosphere. Theoretical modeling neously measured Si I 10827 A of the Hanle and the Paschen–Back effect helped to significantly improve the analysis of polarization measurements in the He I triplet, allowing to directly visualize the magnetic structure of spicules, polar prominences and active regions. Here, I will summarize the results of chromospheric magnetic field measurements using this interesting triplet obtained in the last couple of years and discuss the great potential it has to further uncover the complex structure of the chromosphere and its coupling to the photosphere.  2007 COSPAR. Published by Elsevier Ltd. All rights reserved. Keywords: Sun; Chromosphere; Magnetic field; Stokes polarimetry

1. Introduction ˚ The history of solar observations in the He I 1083.0 A triplet dates back to the year 1934, when the line was discovered by Babcock and Babcock in the solar spectrum. Only a few years later, D’Azambuja and D’Azambuja (1938) obtained already quite remarkable spectroheliograms. The potential of measuring chromospheric magnetic fields using this triplet was emphasized by Harvey and Hall (1971): especially the formation at chromospheric heights only (no contributions from the photosphere in the line wings), the absence of photospheric blends, and the high magnetic sensitivity are of clear advantage compared to other chromospheric lines. Until the availability of spectropolarimetric measurements research concentrated on the investigation of the chromospheric topology and the velocity fields: the first report on steady, long-lasting flows

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Tel.: +49 5556 979 465. E-mail address: [email protected].

with lifetimes from hours up to several days, one of the topics in this review, in this interesting chromospheric line was presented by Lites et al. (1985). The model calculations for the formation of the He I ˚ line by Avrett et al. (1994) showed the strong 10830 A dependence of the line absorption on coronal illumination. This work represented the starting point for the systematic analysis of the He I triplet. The first spectropolarimetric measurements to obtain longitudinal magnetograms and velocity fields in the chromosphere were performed by Penn and Kuhn (1995) at the National Solar Observatory (NSO) Vacuum Tower Telescope (VTT) at Sacramento Peak. Ru¨edi et al. (1996) obtained the first full Stokes vec˚ . The measurements tor measurements in He I 10830 A demonstrated the large potential of this triplet and started the development of dedicated infrared polarimeters. One of the most successful instruments was the Tenerife Infrared Polarimeter at the German Vacuum Tower Telescope (VTT) on Tenerife (TIP, Martı´nez Pillet et al., 1999), which recently was upgraded to increase spatial and spectral resolution as well as to increase the field-of-view to 7700

0273-1177/$30  2007 COSPAR. Published by Elsevier Ltd. All rights reserved. doi:10.1016/j.asr.2007.03.091

A. Lagg / Advances in Space Research 39 (2007) 1734–1740

(TIP-2, Collados et al., 2007). In this review, I would like to present some recent advances in measuring the chromospheric magnetic field, focused on the capacity of Zeeman ˚ multiplet. and Hanle diagnostics in the He I 10830 A This review is structured as follows: in Section 2, I briefly outline the formation of the He I triplet and the methods to analyze the measured Stokes profiles. Section 3 presents a few, selected examples of magnetic field measurements based ˚ . Conclusions and outlook are presented in on He I 10830 A Section 4. 2. The He I triplet ˚ line is a transition from the 3S1 state to The He I 10830 A 3 the P2,1,0 state of the orthohelium. The formation height lies just below the transition region, between 1500 and 2000 km, although this layer may be significantly warped (He I signatures can arise from heights of several tens of Mm in prominences). The population of the ground-level for this transition occurs mainly via ionization of He I due to coronal UV radiation and subsequent recombination of He II (Avrett et al., 1994). Andretta and Jones (1997) stress the importance of the conditions in the lower transition region for the formation of this line, and the strong population imbalance of the orthohelium ground state. The line parameters of the triplet are given in Table 1. The spectral window around the triplet contains ˚ (Lande´ factor a photospheric Si I line at 10827.088 A geff = 1.5), making this region ideally suited to study the coupling between the photosphere and the chromosphere. Whereas the Zeeman effect dominates the linear polarization signature in the Stokes vector for fields larger than 1000 G, the Hanle effect allows the determination of fields in the sub-Gauss to the tens of Gauss regime. The quantum theory to explain the puzzling polarization signature of atomic level polarization and its modification by the Hanle effect was developed by Trujillo Bueno et al. (2002). Since then it has become possible to develop inversion codes for the He I line to obtain maps of the chromospheric magnetic field topology. Socas-Navarro et al. (2004) pointed out the importance of taking into account the Paschen–Back effect for the inversion of Stokes profiles induced by the Zeeman effect. Their work could explain partly the presence of asymmetries in the Stokes V profile even in Milne–Eddington type atmosphere and improved the accuracy of magnetic field and velocity calculations (Sasso et al., 2006). The polynomial approximates provided by Socas-Navarro et al. (2005) facilitate the consideration of the Paschen–Back effect in existing inversion codes Table 1 Line Parameters of the He I triplet ˚) Line Wavelength (A Transition He (Tr1) He (Tr2) He (Tr3)

10829.0911 10830.2501 10830.3397

3

3

2s S1–2p P0 2s3S1–2p3P1 2s3S1–2p3P2

geff

Rel. osc. strength

2.0 1.75 1.25

0.111 0.333 0.556

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based on the linear Zeeman effect theory, such as the non-LTE code of Socas-Navarro et al. (2000) or the Milne–Eddington code of Lagg et al. (2004). Recently, Trujillo Bueno and Asensio Ramos (2007) have demonstrated that the selective emission and selective absorption processes caused by the presence of atomic level polarization may have an important influence on the emergent linear polarization, even for magnetic field strengths as large as 1000 G. Therefore, unless the magnetic field strength is sensibly larger than 1000 G, the modeling of the stokes Q and U profiles should be done by taking into account the joint action of the Zeeman effect, atomic level polarization and the Hanle effect. To this end, Asensio Ramos and Trujillo Bueno (2007) have developed a more general inversion code that takes all such effects properly into account. 3. Examples of He I measurements 3.1. Magnetic expansion and the canopy in the quiet Sun The expansion of magnetic field lines with increasing height is a logical consequence of the decrease in gas pressure. This is the basic principle of the ‘canopy’ models, established by Gabriel (1976) and Giovanelli (1980). In these models, magnetic flux escaping from intergranular lanes expands rapidly with increasing height, leading to the formation of horizontal field lines spanning over supergranular cells. Significant sophistication has been added to these models (e.g. Solanki and Steiner, 1990), with the focus being on the determination of the exact location of the canopy. Numerous measurements give strong evidence for canopy like fields (e.g. the Hanle depolarization in Sr II by Bianda et al. (1998, 1999), Zeeman and Hanle diagnostics in Na D2 by Stenflo et al. (2002), acoustic mapping of the plasma-b = 1 surface by Finsterle et al. (2004)). However, there are also indications that the concept of a magnetic canopy is fundamentally wrong. Zhang and Zhang (2000) presented disk-center, quiet Sun magnetograms in ˚ line and in the chromospheric the photospheric Fe I 5324 A Hb-line obtained with the Huairou magnetograph in the Beijing Astronomical Observatory. They measure little expansion of photospheric magnetic flux with height, which clearly is in contradiction to the classical canopy picture. A similar result was obtained by Schmidt and Tritschler (presented in Peter, 2002) using cospatial images in the C I continuum and the O VI line from the SUMER instrument. Schrijver and Title (2003) questioned the canopy concept: the presence of unresolved, relatively strong inter-network field concentrations of the order of a few M · cm2 destroy the classical, wineglass shaped canopy. Their simulations show that 50% of the coronal field is rooted in the inter-network. This interpretation is supported by high-resolution Gband observations of bright points in the inter-network by Sa´nchez Almeida et al. (2004), obtained with the Swedish Solar Telescope on La Palma. These bright points can be

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interpreted as the inter-network field concentrations determining the vertical structure of the quiet corona. Is the concept of a magnetic canopy spanning over internetwork, weak-field regions fundamentally wrong? Measurements in the He I triplet should be a valuable tool to proof or disproof this concept. The height of formation of this triplet is in the upper chromosphere, a region where the magnetic expansion, if present, should already have happened and the magnetic field vector in the inter-network should be horizontal. First attempts to measure these weak fields are presented in Fig. 1: the Stokes vector represents a 200 · 200 average over the central part of a supergranular cell. The observation was obtained with the new TIP-2 instrument at the German Vacuum Tower Telescope, the heliospheric angle was 60 (l = cosH = 0.5). Clearly, a Stokes V signal is present. An inversion involving Zeeman and a simplified version of Hanle diagnostics (Lagg et al., 2004) reveals magnetic field strengths of <100 G and an inclination angle with respect to the solar surface of 60–120. For this inversion higher weight was given to

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0.001

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10826

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10834

Wavelength [Å]

Fig. 1. Observation of a quiet region: the measured Stokes vector (black) was analyzed using a Milne–Eddington type inversion (Lagg et al., 2004). The best-fit profile is shown in red. The values obtained for the magnetic field vector are jBj < 100 G, inclination to the solar surface c = 60–120. The dashed blue line shows the weighting scheme used for the inversion. (For interpretation of the references to colour in this figure legend, the reader is referred to the web version of this paper.)

the Tr2 and Tr3 of the He I triplet, since Tr1 is blended with a photospheric Ca I line (see weighting scheme indicated by dashed blue line in Fig. 1). Further analysis of this data set, especially the proper treatment of the Hanle effect, is required to confirm the presence of the canopy. Furthermore, the measurement must be repeated over multiple supergranular cells. 3.2. Wave propagation from photosphere to chromosphere The presence of the photospheric Si I line in the close vicinity of the He I triplet makes the spectral region around ˚ an ideal window for the investigation of wave 10830 A propagation from the photosphere to the upper chromosphere. Cospatiality and simultaneity of observations in these two layers, required for such studies, are automatically fulfilled. Using spectropolarimetric data in this spectral window from the TIP instrument (VTT) Centeno et al. (2006) (see Fig. 2) were able to calculate the phase relation between the velocity oscillations in the Si I and the He I line. They interpreted the phase differences by vertically propagating, slow magneto-acoustic waves in a stratified magnetized atmosphere taking into account the damping of the temperature fluctuations due to radiative losses according to Newton’s cooling law. They conclude that the dominant 3-min power in the chromosphere directly comes from the photosphere by means of linear wave propagation and that no non-linear interaction of 5-min wave modes is required. The saw-tooth like structure of the chromospheric waves is nicely explained by the amplitude increase due to the decreasing density of the upward propagating waves, leading to the development of shock waves at chromospheric heights. Bloomfield et al. (2007) extended this work by taking into account non-vertical wave propagation along the magnetic field lines. Their analysis confirms the conclusions by Centeno et al. (2006), and additionally provides a possible explanation for the running penumbral waves (RPWs) in the chromosphere: these waves, already observed by Zirin and Stein in 1972, propagate outward from the umbra with speeds between 15 and 20 km s1. RPWs have been studied extensively (e.g. by Christopoulou et al., 2001), but the origin of these moving disturbances and their relation to other phenomena occurring within sunspots remained unclear. Bloomfield et al. (2007) explain RPWs by the increase of the travel distance of the waves caused by the increase of the magnetic field inclination with larger distance from the umbra. If the chromospheric waves are indeed excited in the photosphere this increasing path length would lead to the observed, apparent outward motion of the chromospheric velocity signal (as already proposed by Rouppe van der Voort et al., 2003). 3.3. The 3-D structure of a sunspot Modern inversion techniques, like SPINOR (Frutiger, 2000) or SIR (Ruiz Cobo, 1998), are able to determine

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˚ line (top) and the chromospheric He I 10830 A ˚ line (bottom) inside the umbra of a sunspot (May 9, Fig. 2. Velocity maps for the photospheric Si I 10827 A 2001). Clearly visible are the 5-min oscillations in the photosphere and the 3-min oscillations in the chromosphere (taken from Centeno et al., 2006).

the parameters of the solar atmosphere to a high degree of sophistication. Multi-component, height-dependent models can easily be computed. However, the quality of the determined atmospheric parameters is naturally limited by the information contained in the Stokes spectra: multiple line spectropolarimetry, successfully performed, e.g. with the THEMIS telescope (e.g. Molodij and Rayrole, 2003), provides the best information on the stratification of the solar ˚ does not atmosphere. The Si I/He I line pair at 10830 A exhibit a continuous height sampling of the solar atmosphere, since it only provides the atmospheric parameters in two relatively thin layers in the photosphere and the chromosphere, respectively. However, the large separation of these layers and the intrinsic cospatiality of the observations allow for a detailed analysis of the three-dimensional structure of the magnetic field topology of active regions. Such an analysis was presented by Orozco Suarez et al. (2005) who investigated the gradient of the magnetic field in a sunspot (NOAA 10130, September 28 2002, TIP/ VTT). Similar to the work of Eibe et al. (2002), who determined the gradient using a response function based method applied to the Na I D1 line, Orozco Suarez et al. calculated the photospheric gradient solely from inversions of the ˚ line. The azimuthal average of photospheric Si I 10827 A this gradient is plotted in Fig. 3 as dark gray shaded area. The gradients in umbra and penumbra are in good agreement with the inversions by Westendorp Plaza et al. (2001). In addition to the response function based inversion of the Si I line, Orozco Suarez et al. determined the magnetic field at chromospheric heights using a Milne–Eddington type inversion code (Lagg et al., 2004). Assuming a difference in formation height of 2000 km between the Si I and the He I line the so computed magnetic field gradient (light gray area in Fig. 3) is on average significantly lower, especially in the penumbra. In a future work, the height separation between the photospheric and the chromo-

Fig. 3. Magnetic field gradients determined from the LTE-inversion of the ˚ line (black dotted line + dark gray area) and from the Si I 10827 A difference between the photospheric and the chromospheric layer (red dotted line and light gray area). The vertical dashed lines indicate the boundaries between umbra, penumbra and surrounding quiet Sun, respectively. The scatter indicated by the shading results from the azimuthal assuming an elliptical shape of the sunspot. (For interpretation of the references to colour in this figure legend, the reader is referred to the web version of this paper.)

spheric layer might be constrained better by applying full non-LTE inversions of the He I line. However, to retrieve a complete picture on the 3-D topology of a sunspot multi-line investigations sampling the heights between these two lines are mandatory. Socas-Navarro (2005) performed such a multi-line analysis using the SPINOR1 spectropolarimeter at the Dunn solar Telescope. From their three-dimensional map of the magnetic field vector between 1 This instrument shares the same abbreviation as the SPINOR inversion code mentioned above.

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3.5. Multi component downflows At chromospheric heights the magnetic field is believed to fill the entire volume, i.e. the magnetic filling factor is

20000 18000 16000 14000 12000 10000 8000 0.003 0.002 0.001 0.000 –0.001 –0.002 –0.003 0.003 0.002 0.001 0.000 –0.001 –0.002 –0.003 0.02

0.0 1.4 1.2 1.0 0.8 0.6 0.4 0.2 0.0 1.4 1.2 1.0 0.8 0.6 0.4 0.2 0.0

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In a pioneering work Bommier et al. (1994) determined the average magnetic field strength in prominences from the linear polarization signature in the He I D3 and Ha lines. They obtain field strengths of 7.5 ± 1.2 G and an average inclination angle with respect to the solar vertical of 60. Other investigations, nicely summarized in the introduction of Paletou et al. (2001), lead to similar results for magnetic field strength and orientation in prominences. Lin et al. (1998, 2001) were the first to successfully measure ˚ polarimthe magnetic field in a filament using He I 10830 A etry. However, their classical formalism failed to reproduce the puzzling signature of the reversal of the Stokes Q component between the main component (Tr2 and Tr3) and the blue component (Tr1, see Table 1). The quantum mechanical approach by Trujillo Bueno et al. (2002) solved this puzzling signature of scattering polarization and its modi˚ line fication due to the Hanle effect in the He I 10830 A by taking into account atomic polarization of the lower level of this triplet. Based on this work, Merenda et al. (2006) were able to determine the magnetic field vector in a polar crown prominence and found significant difference to the expected values: the magnetic field strength is higher (30 G instead of expected 7.5 G) and the inclination to the solar vertical is much lower (24, expected 60). They also observe a rotation of the magnetic field in the central part of the prominence. It remains to be checked whether the prominence analyzed by Merenda et al. (2006) was a rare exception or the theoretical models of prominences have to be refined. Similar values for magnetic fields in prominences (30–45 G) were obtained by Paletou et al. (2001) using multi-line polarimetry at the THEMIS telescope. Trujillo Bueno et al. (2005) presented the first direct empirical demonstration of a significant magnetization of ˚ line and spicules using observations in the He I 10830 A theoretical modeling of the Hanle and the Zeeman effect. They determined the magnetic field strength in spicules at a height of 2000 km to be 10 G, which they believe is a typical value for quiet Sun spicular material. Similar values for the average magnetic field strength in spicules were obtained by Lo´pez Ariste and Casini (2005) by analyzing full Stokes vector ASP observations in the He I D3 line. They conclude that the magnetic field in spicules is aligned with the visible structures and measure field strengths of up to 40 G. Combining ground-based measurements in He I ˚ multiplet bear a great potential D3 line and the 10830 A to uncover the magnetic topology of spicules in the near future.

Q/IC

3.4. Magnetic field in prominences and spicules

unity. The expansion of the magnetic field structure might lead to the conclusion that any magnetic fine structure present in lower layers gets lost in the chromosphere, resulting in a smooth, uniform magnetic field topology. High-resolution Ha observations already proof that there is significant fine structure in the chromosphere. Also, highest spatial resolution observations currently available in the ˚ line (0.700 , TIP2/VTT) show unresolved fine He I 10830 A structure. Lagg et al. (2007) present an example of such an unresolved fine structure in the chromosphere: close to a pore (NOAA 9451, May 13 2001). They observe Stokes profiles showing two clearly distinct velocity components (see Fig. 4). Such two-component profiles were already observed previously (e.g. Schmidt et al., 2000), but now the magnetic structure of the two components could be determined independently. In contrary to the assumption of Schmidt et al. (2000), Aznar Cuadrado et al. (2005, 2007) find that multi-component profiles are very common in active regions. Using a Milne–Eddington inversion combined with a robust genetic algorithm (PIKAIA) Lagg et al.

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0 and 1800 km they were able to determine the twist and the torsion of the magnetic field in a sunspot.

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Fig. 4. Stokes profiles showing two distinct magnetic components in the close vicinity of a pore (from Lagg et al., 2007). The observed profiles are shown in black, the fits in red (two independent magnetic components, Bslow = 725 G, cslow = 30 and Bfast = 1194 G, cfast = 68) and green (2 magnetic components with coupled magnetic field Bslow = Bfast = 642 G, cslow = cfast = 61). The two components are shifted by 28 km s1 relative to each other. The orange line shows a fit assuming one broad absorbing component superposed with a central emission component (fitness f = 1.50). The fit involving two independent magnetic components best reproduces the observed Stokes vector. The dashed, blue line indicates the weighting scheme used for the inversion, the vertical, dotted lines show the central wavelengths of the three components of the He I triplet. (For interpretation of the references to colour in this figure legend, the reader is referred to the web version of this paper.)

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found that it is not possible to obtain a satisfactory fit to the observed profiles assuming the same magnetic field topology for both velocity components; a different magnetic field orientation for the two components is mandatory. Two interpretations for this observation are possible: (1) the magnetic field structure in the chromosphere is highly filamentary, similar to the uncombed structure of the penumbra in the photosphere, or (2) the He I absorption results from two optically thin magnetic clouds at different heights along the line-of-sight path. Further observations under various viewing geometries and higher spatial resolution observations as expected with the new GREGOR telescope will help to clarify this issue. 4. Conclusions and outlook The potential to unveil physical processes responsible for the coupling between the photosphere and the corona by observations in this interesting spectral window is huge. With the analysis tools which have been developed over the last years and their continuous improvements vector magnetic field maps will become a standard product within the next years. Additionally, improved instrumentation, like the TIP-2 instrument mounted behind the new 1.5 m GREGOR telescope on Tenerife or the Spectro-Polarimeter for Infrared and Optical Regions (SPINOR,2 Socas-Navarro et al., 2006) will increase the spatial resolution as well as the signal to noise ratio, allowing to better resolve the fine structure of the chromosphere and to measure weaker magnetic fields. The detailed structure of the chromospheric magnetism, one important ingredient for the solution of the coronal heating problem, will become accessible. References Andretta, V., Jones, H.P. On the role of the solar corona and transition region in the excitation of the spectrum of neutral helium. ApJ 489, 375-+, 1997. Asensio Ramos, A., Trujillo Bueno, J. A user-friendly code to diagnose chromospheric plasmas, in: Heinzel, P., Dorotovic, I.R.R. (Eds.), The Physics of Chromospheric Plasmas, Coimbra, Portugal (October 9–13, 2006), Vol. 368 of ASP Conf. Series. Astronomical Society of the Pacific, pp. 163–170, 2007. Avrett, E.H., Fontenla, J.M., Loeser, R. Formation of the Solar 10830 Angstrom line, in: Rabin, D.M. (Ed.), IAU Symp. 154: Infrared Solar Physics. Kluwer Academic Publishers, Dordrecht, pp. 35–47, 1994. Aznar Cuadrado, R., Solanki, S.K., Lagg, A. Supersonic flows in the solar chromosphere are very common, in: Innes, D.E., Lagg, A., Solanki, S.K., Danesy, D. (Eds.), Chromospheric and Coronal Magnetic Fields, Vol. ESA SP-596. ESA Publication Division, November 2005. Aznar Cuadrado, R., Solanki, S.K., Lagg, A. Velocity distribution of chromospheric downflows, in: Proceedings of Modern Solar Facilities – Advanced Solar Science Workshop, Go¨ttingen, Germany, 2007. Babcock, H.D., Babcock, H.W. Some new features of the solar spectrum. Publ. Astron. Soc. Pac. 46, 132–133, 1934.

2

The abbreviation SPINOR originally has been used for the inversion code ‘Stokes-Profiles-INversion-O-Routines’, developed by Frutiger (2000).

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