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Haosheng Lin Institute for Astronomy, University of Hawaii

Current Status of Solar Magnetism Research Large Off-Axis Reflecting Coronagraph and Advanced Solar Instrumentation. Haosheng Lin Institute for Astronomy, University of Hawaii. Outline of This Talk. Overview of Solar Magnetism and Unresolved Problems—solar dynamo Sunspot magnetic fields

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Haosheng Lin Institute for Astronomy, University of Hawaii

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  1. Current Status of Solar Magnetism ResearchLarge Off-Axis Reflecting Coronagraphand Advanced Solar Instrumentation Haosheng Lin Institute for Astronomy, University of Hawaii

  2. Outline of This Talk • Overview of Solar Magnetism and Unresolved Problems—solar dynamo • Sunspot magnetic fields • Quiet Sun magnetic fields • Coronal magnetic fields • Why do we need large aperturSe solar telescope • Large Aperture Coronagraph • Why do we need large aperture? • Why do we need a coronagraph? • Why Off-Axis coronagraph? • Ongoing and Expected Off-Axis Coronagraph Projects • Advanced Instrumentation • Clever multiplexing to achieve high spatial, spectral, and temporal resolution and sensitivity with simultaneous 2-D spatial coverage • What can we do with a coronagraph during the night?

  3. Solar Magnetism—An Overview Sunspots and active regions are manifestation of a global-scale solar dynamo operating at the base of the solar convection zone. The existence of magnetic fields makes the Sun an interesting star to study… ‘Quiet’ Sun ‘Quiet’ Sun Line-of-sight Magnetic Flux Continuum Intensity

  4. Solar Cycle Sunspot Cycle The number of sunspots on the surface of the Sun follows a 11-year cycle. Butterfly diagram Sunspots appear at higher latitude at the beginning of the solar cycle, and migrate toward the equator, as the cycle evolve

  5. Magnetic Butterfly Diagram

  6. S 11 years later… N Magnetic Field Configuration of the Sun During Different Phase of the Solar Cycle • Solar Minimum • Dipole Magnetic Field • No Sunspot • Solar Maximum • ToroidalMagnetic Field • Many Sunspots ~5 years later… The magnetic field configuration of the Sun evolves with a 22 year cycle. 22 years later…

  7. The solar dynamo is believed to be generated at the base of the convection, where the rotation rate changes abruptly.

  8. MHD instability causes the flux tubes to raise to the surface—sunspots. ‘Quiet’ Sun Continuum Intensity

  9. Current Solar Magnetism Research Emphasis—ATST Science Goals • 2.2 Magnetic Field Generation and the Solar Cycle 7 • 2.3 Magnetic Flux Emergence: Active Region Emergence and Evolution 9 • 2.3.1 Surface Manifestation of Subphotospheric Processes 10 • 2.4 Small-Scale Magnetic Flux Concentrations 11 • 2.5 Magnetoconvection and Flux Transport 15 • 2.5.1 Sunspots: An Example of Magnetoconvection 16 • 2.5.2 Formation of sunspots 19 • 2.5.3 Radiative Flux balance in sunspots. 19 • 2.5.4 Filamentary Structure of Sunspot Penumbrae 20 • 2.6 Structure and Dynamics of the Chromosphere 21 • 2.7 Coronal Structure and Heating 23 • 2.7.1 Coronal Structures in Three-Dimensions 25 • 2.7.2 Electric Current Systems 26 • 2.7.3 Coronal Magnetic Fields: The Need for Ground-Based Coronal Studies 27 • 2.7.4 New Spectral Diagnostics 30 • 2.7.5 Wide-Field Coronal Photometry 31 • 2.8 Flux Expulsion: Flares and Mass Ejections 31 • 2.8.1 High Resolution, High Cadence Studies of Solar Flares 33 • 2.8.2 Structure and Evolution of Magnetic Fields Associated with Flares and CMEs 34 • 2.8.3 CMEs: The need for prominence studies 35

  10. Sunspot: The Oldest Mystery of Solar Physics Ancient Chinese noted the existence of sunspots more than two thousand years ago. Galileo observed sunspots through the first telescope he made…but sunspot phenomenon remains one of the oldest mystery of solar physics. We know a lot about sunspots observationally, but we still do not have a self consistent model of sunspot today. • Why are sunspots dark against the photosphere? • Why is there a sharp boundary between the photosphere and the penumbra? • Why is there a sharp boundary between the penumbra umbra? • Why is there an Evershed flow? • … • … Swedish Vacuum Tower Telescope Image

  11. Dynamic of Sunspot Swedish Vacuum Tower Telescope Movie

  12. What is a Sunspot? • G. E. Hale (1908) first demonstrated that sunspots are regions of strong (2,000 ~ 3,000 Gauss) magnetic fields… • Zeeman Effect in Sunspot • Zeeman splitting B is proportional to 2. Therefore, it is preferable to perform magnetic field observation in the infrared wavelength: • B = 4.67 10-7g2B • where g is the Lande factor The FeI 1564.8 nm (g =3) and FeI 1564.3 (g=1.53) nm line pair are the most sensitive magnetic field diagnostics in the visible and near-IR wavelengths. Visible lines (such as FeI 630.2 nm line (g=2.5) offers lower magnetic sensitivity, but higher spatial resolution.

  13. Why is Sunspot Dark? Zero-th Order Explanation: Given the fact that there sunspots are regions with strong magnetic fields, and that sunspots are long-lived, stable features, magnetohydrostatic equilibrium must exist between the magnetized plasma of the sunspots, and the non (or weakly)-magnetized plasma outside of the sunspot. The presence of the magnetic pressure (B2/8) requires the sunspot plasma to be cooler in temperature… Pphoto(z) = Pspot(r,z) + Bz2(r,z)/8 + Fc(r,z)/8 Where, Fc(r,z) = 2 ∫Bz(r’, z) (Br(r’, z) /z) dr’ Pphoto(z) = n(z) k Tphoto(z) Pspot(z) = n(z) k Tspot(z) A carton model of sunspot Tspot ~ 3,500 K B ~ 3,000 Gauss Tphoto ~ 6,000 K Bphoto ~ 0 Gauss So, sunspots are dark because of the presence of strong magnetic fields. But why is there a penumbra, and an umbra?

  14. Is Sunspot in MHS Equilibrium? Kopp & Rabin found different slope of the B2 vs. dT curve for different part of sunspot with IR FeI 1564.8 nm line measurements. Martinez et al (1993) found a linear relation between B2 and dT with magnetic field measurements in the visible wavelength (FeI 630.2 nm line).

  15. Although the idea about magnetohydrostatic equilibrium between the magnetized sunspot atmosphere and its surrounding is plausible, observational verification (or dismissal) has not been demonstrated yet. Furthermore, given the complicated structures of the sunspots, we should not expect a simple relationship between B and T. Magnetic field measurements have not achieved the same spatial resolution as those demonstrated from the new Swedish Vacuum Tower Solar Telescope yet. Obviously, better magnetic field data are needed… ‘Good’ resolution visible vector magnetic field measurement

  16. Polarized Spectra of Sunspot • Solar radiation in general is not polarized, except when observing near or off the solar limb where scattering polarizations are significant. However, magnetic fields modify the absorption and emission coefficients of Zeeman sensitive spectral lines and produce polarized spectral radiation. Therefore, the polarization states of solar spectral radiation carry information about the vector magnetic fieldconfiguration of solar magnetic fields.

  17. Measurement of Flux vs. Field Strength B2 B1 B3 • Longitudinal magnetograph (SOHO MDI, BBSO magnetogram) measures the magnetic fluxwithin the resolution element along the line-of-sight… • B = ∫sB · da • Magnetic field strength |B| can only be measured by direct (spectroscopic) measurement of the Zeeman splitting… • Minimum measurable |B| depends on the doppler width of the spectral line. • Structures within the resolution elements are ignored. ‘pixel’ of observation

  18. New High-Resolution IR Measurements of Sunspot Magnetic Fields Magnetic Field Strength B

  19. OH Equivalent Width OH equivalent Width vs temperature is independent of the sunspot size  The formation of OH molecules depends only on the temperature of the sunspot, not magnetic fields?

  20. B vs. OH • OH molecules form only in the umbrae of the sunspots… • Dissociation energy 4.3 eV • Slopes of B-OH curves are different for different size sunspots…

  21. B vs Tc • B-T curve is a strong function of sunspot size… • Phase transition in sunspot? • Theoretical calculation predicted that about 20% of Hydrogen are in molecular form (H2)… • Molecules have internal degree of freedom (rotation and vibration). • The formation of molecules introduce a discontinuous change in the heat capacity of the sunspot plasma  first-order phase transition? • What’s the thermodynamic effect of molecules in sunspot atmosphere?

  22. Quiet Sun Magnetic Fields • ‘Quiet’ Sun is not really magnetic-field-free. Weak field (<1,000 G) features can be seen everywhere… • Quiet Sun magnetic fields may be generated by a small-scale surface dynamo… ‘Quiet’ Sun ‘Quiet’ Sun SOHO MDI ‘White Light’ Image SOHO MDI Magnetogram

  23. Example of Weak Field Stokes V Spectra • Precision IR spectropolarimetry allows for determination of |B|down to ~ 150 G. • Sensitivity to B ~ 5 x 1015 G·cm2 • Sensitivity to B is limited only by photon noise!

  24. Quiet Sun Magnetic Fields • Small-scale, weak magnetic fields with mixed polarities can be found every where on the surface of the Sun. • These fields are mostly cospatial with the intergranular lane, and evolve with a time scale similar to that of solar granulation • <B> ≈ 500 G, equipartition with turbulent motion of the solar granulation Intensity Kilo-gauss component in equipartition with the gas pressure of the photosphere Magnetogram

  25. Pending Questions about the Weak Fields • How are the weak field component generated? • Local dynamo associated with the convective motions? • Recycling of the magnetic fluxes generated by the global scale dynamo? • Does ‘flux tube’ exist? • If the quiet Sun magnetic fields are generated by a local solar dynamo, then we expect the magnetic fields to exist in every spatial scale and in equipartition with the turbulent gas pressure. This ‘turbulent magnetic fields’ will evolve with the same time scale as the turbulent flows. • In order to answer the pending questions concerning the quiet Sun magnetic fields, we need to observe the quiet Sun magnetic field with the highest spatial resolution and magnetic field sensitivity attainable, and with the highest temporal resolution possible and simultaneous 2-dimensional coverage—this statement is in fact true for almost every aspect of solar magnetism research now…

  26. Coronal Magnetic Fields Coronal magnetic field is something of a dark energy problem for solar physics in that we know it permeates the corona and controls its static and dynamic behavior, yet we are unable to usefully measure it...Because of the high temperature (T ~ 106 K) and low field strength (B ~ 10 G) condition of the solar corona,measurement of coronal magnetic field is one of the most challenging task of observational solar astronomy. Without direct measurement, coronal intensity loop structures are thought to be proxies of coronal magnetic field. But, observational verification is need! This talk is about how we can measure coronal magnetic field today… SOHO/EIT Fe XVI 284 Å

  27. Coronal Magnetic Fields and Coronal Mass Ejections and Flares The energy released in coronal mass ejects (CMEs) and flares are thought to be stored in the coronal magnetic fields… SOHO/LASCO images of coronal mass ejection TRACE image of post-flare loops

  28. Tools for Coronal Magnetic Field Diagnostics • Faraday rotation of astronomical radio source (Patzold et al., 1987) • Faraday rotation of polarized solar radio ration (Alissandrakis and Drago, 1995) • Gyrosynchrotron radiation magnetometry (Gary and Hurford, 1994, coronal B at the ‘base’ of the corona, on the solar disk) • UV (E1) Hanle effect of O VI 103.2 nm (Sahal-Brechot et al., 1986) • —depolarization and rotation of linear polarization… • Extrapolation from photospheric magnetic field measurements • M1 (magnetic dipole) Hanle and Zeeman effect polarimetry • Stokes V gives strength of B • Stokes Q and U yield orientation of B projected in the plane of the sky with a 90 ambiguity MHD simulation of pre-CME coronal magnetic field structure, Roussev et al. 2003

  29. Difficulties of Coronal Magnetometry • Due to the high temperature (106 K)and low magnetic field strength (B ~ 10 G) of the corona: • V ~ few × 10-4IL  Need 108 photons per measurements! • Low photon flux (10–5 of disk center intensity)from the solar corona • Large scattered background • Q, U ~ 1 × 10-1IL >> V •  Just a very small linear-to-circular polarization crosstalk is sufficient to mask the weak Stokes V profiles.

  30. History of CEL Coronal B Observations • Early Attempt • Harvey, 1969: Fe XIV 530 nm Stokes V magnetometry • No detection. • Linear Polarization Maps • Mickey, 1973: Fe XIV 530 nm • Querfeld, and Smartt, 1984: Fe XIII 1075 nm • Arnaud & Newkird, 1983: Fe XIII 1075 nm •  Successfully obtained maps of the orientation of coronal magnetic fields. • Recent Efforts • Kuhn, 1995: Fe XIII 1075 nm Stokes Vspectropolarimetry • No detection. • Lin, Penn, & Tomczyk, 2000: Fe XIII 1075 nm Stokes V spectropolarimetry •  First Definitive detection of line-of-sight coronal magnetic field! Linear polarization map of the Fe XIII 1075 nm line. Habbal et al. 2001, Arnauld 1983

  31. CEL Polarimetry—Incoherent Magnetic Resonant Scattering • Physical Process: • Resonant scattering of anisotropic • photospheric radiation by atoms and ions in the corona in the presence of a magnetic field. • Linear Polarization • – Orientation of CEL linear polarization maps the orientation of magnetic field projected in the plane-of-sky (POS) • – Magnetic field orientation subject to 90 degree ambiguity (Van Vleck Effect). • Circular Polarization • – Circular polarization of CEL is proportional to the strength of line-of-sight magnetic field • – The magnetograph formula is modified by an alignment factor that depends on the inclination angle between B and the local vertical direction, and the anisotropy of the incident radiation field.

  32. Physics of Resonance Scattering Polarization wB wB A A

  33. Stokes Parameters of CEL • Diagnostics of Coronal B • U/Q = tan 2 • V ~ cos  vB (dI/d) • P = (Q2+U2)-1/2/I ~ f (,M) sin2  • Q and U vanish at van Vleck angle vv = 54.7º • Q and U change sign for  > vv

  34. Example of Fe XIII 1075 nm Coronal Emission Line Spectropolarimetry Magnetogram of target active region observed on disk Target of Observation 1999/10/25 Weak Stokes V signal in the FeXIII 1074.7 nm line can be detected!

  35. First Definitive Coronal Stokes V Measurement • NSO/SP Evans Solar Facilities 40 cm coronagraph • 240 arcsec2 FOV (summed over the entire length of the slit). • 2560 seconds (44 minutes) integration time (Q & V). • Careful telescope and instrumental polarization cross-talk control • Coronal magnetic field can be measured with Zeeman effect diagnostics!

  36. Can We Make 2-D Coronal Magnetic Field Maps? IfA Coronal B Initiatives • While Lin et al. (2000) demonstrated the feasibility of using CEL polarimetry to measure the strength of coronal magnetic fields, useful measurements require 2-dimensional spatial coverage. • Given the long integration time required to obtain one measurement, can we make 2-D coronal magnetic field maps? • To this goal, we initiated a new effort to establish the capability to make regular 2-dimensional maps of both longitudinal magnetic field strength and the orientation of the magnetic field projected in the plane of the sky. • IfA effort includes: • Construction of a 50 cm aperture off-axis mirror coronagraph—SOLARC • Construction of an Optical Fiber-bundle Imaging Spectropolarimeter (OFIS): The polarized spectra of a extended 2-D target can be obtained simultaneously without scanning of the spectrograph slit.

  37. SOLARC: Off-Axis Mirror Coronagraph • PI—Jeff Kuhn (IfA) • 50 cm aperture off-axis gregorian telescope • No secondary mirror and spider structure in the optical path for coronagraphic performance Secondary mirror Prime focus inverse occulter/field stop Re-imaging lens LCVR Polarimeter Input array of fiber optics bundle SOLARC and its dome on the summit of Haleakala, Maui. Primary mirror Optical Configuration of SOLARC and OFIS

  38. OFIS: A True Imaging Spectropolarimeter Echelle Grating Camera Lens Collimator NICMOS3 IR camera Fiber Bundle • NICMOS3 IR Camera • 16  8 => 2  64 optical fiber-bundle • 160  308 mm, 79 lines/mm echelle grating with 63.5 blaze angle • f = 800 mm,  = 150 mm (F/5.3) collimator and camera lens The coherent optical fiber-bundle rearrange the 2-dimensional image sampled by the 16 × 8 input array to two linear array (2 × 64 ). The two linear arrays act as the slits of the spectrograph, thus allowing for the simultaneous recording of the spectra from all the field points in the 2-D image plane.

  39. Sample CEL Spectra from OFIS One 64-fiber column illuminated 16 × 4 pixels area coverage Two 64-fiber columns illuminated 16 × 8 pixels area coverage

  40. 2004/04/06 Observation • Full Stokes vector observations were obtained on April 6, 2004 on active region NOAA 0581 during its west limb transit. • Corona activity is low compared with the 1999 observations! • Stokes I, Q, U, & V Observation: • 20arcsec/pixel resolution • Telescope pointing @ • Radius Vector 0.25 R • Position Angle (Geocentric): 260°. • 70 minutes integration on V • 15 minutes integration on Q & U • Stokes I, Q & U Scan: • RV = 0.25 R • From PAG 250° to 270° • Five 5° steps Fe X 171Å image of the solar corona at approximately the time of SOLARC/OFIS observation from EIT 195 A.

  41. Full Stokes Spectra of CEL • CEL intensity falls off as a function of height h, • Linear polarization increases with h, • As expected from theory • Spectral characteristics of Stokes V similar to I (and Q, U), • Spatial variation of V resembles that of U, •  Strong linear to circular polarization crosstalk! Full Stokes pectra obtained above NOAA 0581. The display ranges for I, Q, U, and V are -0.5 IC to 0.5 IC, -0.05 IC to 0.05 IC, -0.05 IC to 0.05 IC, and -0.005 IC to 0.005 IC, respectively.

  42. Polarization Crosstalk Correction The crosstalk contaminated circular polarization V’ ()can be expressed by V’ = V + a ·Q + b ·U = V + ·I, where V () is the uncontaminated circular polarization signal, a and b are the Q-V and U-V crosstalk coefficients, respectively, and  is an ‘apparent’ I-V crosstalk coefficient. Using a least squares algorithm minimizing 2 =  (V’ – V - ·I), we can derived  assuming I ·V=0 due to the antisymmetric property of V,  = I ·V’ /  I 2. Also, since in weak-field approximation, V = B·dI/d, Row Stokes V and crosstalk-corrected V. The image is rearranged such that the each 8-fiber strip in the vertical direction corresponds to a 8-fiber column in the north-south direction. The first 8 rows (0-7) correspond the column closest to the solar limb. The weak antisymmetric V profiles can be seen in the first two north-south columns (fiber 0 to 16) in the crosstalk-corrected V image. the observed circular polarization can be written as V’ () =  ·I () + B ·dI () /d =  ·I (+ B/), Thus, B can be directly measured by comparison with the shift of V with respect to I in the spectral direction.

  43. Line-of-Sight Magnetic Fields B B Samples of measured and fitted Stokes I and V spectra of the 10  4 (200”  80”) pixel region closest to the solar limb. The errors of the magnetic fields are 1 sigma error. Geocentric north is up, and east is left. The longitudinal field reverses sign around h=0.17 R!

  44. Radial Variations of B and Comparison with Model Calculations • Solid line with error bar IR data • The dotted line Abbett et al. (2003) near-limb magnetic model scaled to an active region with 1000G longitudinal field strength at the photosphere. • * with error bars Global Ledvina et al. (2004) B model (rms field evaluated along averaged horizontal sight path). The upper error bars show the maximum field at given horizontal level, the lower error flag shows the standard deviation of the model B and the plotted symbols show the mean rms B at the given horizontal level.

  45. What Light’s Up Some Field Lines? – Is CEL Intensity Correlated with Magnetic Field Strength? • Magnetic fields fill the entire volume of the corona. However, intensity images of coronal emission lines always show highly distinctive loop structures. • Why some of the magnetic field lines are filled with high density highly ionized atoms, while the adjacent dark regions are not? • Why does the corona appear different in different CEL images? • Are bright coronal loops actually representative of higher magnetic field strength regions? • Do they actually trace the magnetic field lines? • We found NO correlation between |B| and Iline • bright CEL emission does not necessary imply stronger magnetic fields.

  46. Transverse Magnetic Field Orientation 1 • Does intensity loops track magnetic field lines? • Yes, and No? • See boxes 1 and 2 • Degree of polarization decreases as a function of height=> higher anisotropy and less collisional depolarization. • Van Vleck effect in box 2? • Need More Work! 2

  47. ‘Vector’ Coronal Magnetogram Transverse field orientation Longitudinal Field Strength Contour plot of the line-of-sight magnetogram over-plotted on the EIT FeXVI 284 A image. The contours are 5G, 3G, and 1G.

  48. Summary • We have successfully obtained the first coronal magnetogram, with measurements of both the longitudinal magnetic field strength and orientation of the magnetic field projected in the plane of the sky. The magnetic field sensitivity is ~ 1 G near the limb with approximately 1 hr integration with a 20”  20” spatial resolution. • We observed a radial fall-off of B qualitatively similar to that predicted by some numerical models. • We observed a non-radial magnetic field configuration similar to that implied by the EIT image. However, it is still not clear if the loop structures in the EIT image actually follow the magnetic field lines we measure in these FeXIII data. More studies are needed. • We find no correlation between the brightest emission structures and the strongest longitudinal magnetic fields • Need vary large coronagraph to improve on sensitivity and resolution!

  49. Current Status of Coronal B Experiment… B B The Sun is now in the bottom of the solar minimum…there isn’t much to observed, especially with the 50 cm aperture of SOLARC… Are we actually seeing the reversal of B because we are looking at the loop structure edge-on in the 2004 data? We (Lin and Petrie, 2005) have started an effort to model the coronal magnetic field structure by extrapolation of the observed photospheric magnetic field…If we can match the height of the longitudinal field reversal point, then it will be a triumph of this effort.

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