II. Science Objectives

A. Sunrise and Small-Scale Physical Processes

Both observations and theory/modeling point to small-scale processes — those acting at size scales comparable to (or smaller than) a photospheric scale height — as the agents that control the structure, dynamics, and energetics of the solar atmosphere at larger scales. The small-scale processes are the source of solar variability and, ultimately, solar influences on the Earth. To illustrate the need for the high angular resolution and quantitative diagnostics afforded by Sunrise, we briefly explore only a few examples of important small-scale processes:

  • The mechanism by which the magnetic field is concentrated into intense “flux tubes” in the photosphere,
  • The “cancellation” of magnetic flux, and
  • The generation of MHD disturbances in the photosphere, via interaction of flux tubes with convective motions, and their subsequent propagation to and dissipation within the upper layers.

At present, these processes are not well understood, largely because of the lack of continuous, quantitative measurements of the magnetic field and related atmospheric structure on the relevant size for dynamics: the density scale height in the photosphere. Observations needed to address these and other important issues will involve quantitative measurements of the vector magnetic field, Doppler velocities, and thermal structure. Even the best quantitative field observations to date have angular resolution roughly one order of magnitude worse than that needed to spatially resolve a typical elemental flux tube.

Theoretical mechanisms have been proposed for each of the three examples, but conclusive evidence in support of any specific mechanism has generally eluded observational investigation. Convective intensification (Parker 1978) is a popular explanation for the coalescence of photospheric fields into tiny kiloGauss flux tubes. Figure 1 presents one frame of a numerical simulation of convective intensification. These simulations suggest kilo-Gauss flux bundles form with diameters of order 100 km, and with dynamically important internal structure at least as small as ~30km.

Figure 1:Convective intensification of magnetic field in a numerical simulation (Grossmann-Doerth, Schüssler & Steiner 1998). Shown is a vertical cut through the photosphere and the upper levels of the convection zone, with contour lines outlining the magnetic field. The nearly horizontal curve indicates continuum optical depth unity at 500 nm. A magnetic element with kG field strength has formed through the combined effects of magnetic flux advection by horizontal flow, radiative cooling, and suppression of convection by the growing magnetic field.

Reconnection followed by either resubmergence or expulsioncould explain the apparent cancellation of unrelated flux bundles when they are brought together by larger scale flows (such as the supergranular flow). With very high resolution time series of quantitative vector magnetic field measurements and simultaneous imaging and spectroscopy of chromospheric structure, one may be able to ascertain heights at which reconnection takes place. The relative fraction of flux that undergoes resubmergence as compared to that which is expelled upwards as buoyant U-loops will then be determined. Figure 2 shows a region of apparent flux cancellation within an active region, where full Stokes profile measurements indicate apparent contact of opposite polarity flux.

Figure 2:Advanced Stokes Polarimeter observations of the vector magnetic field in a newly emerging active region reveal a contact region of opposite polarities between the two large, opposite-polarity pores. Opposite polarities of nearly vertical magnetic fields occur within the resolution size of about 1”, as seen in the expanded perspective view in the top plane where arrows represent the direction of the field vector.

The vector field analysis indicates that the field is concave upwards (“U-loop”), but within the contact region there are no extensive areas of horizontal field like those encountered in typical (“-loop”) emerging geometry. Apparently an angular resolution much better than the 1 arc second resolution of these observations will be necessary to explore the physics of the contact region. In the case illustrated here, one would either find smaller-scale, rapidly evolving regions with -loop geometry indicating local reconnection, or persistent U-loop geometry with horizontal fields indicating the buoyant rise of U-loops through the photosphere.

Simulations also reveal that granular convection may induce propagating disturbances in flux tubes (Steiner et al. 1998). This interaction has long been suspected of producing heating of the upper solar atmosphere via the generation and upward propagation of MHD waves, but the extremely high angular resolution needed to detect these very small events in the photosphere has eluded direct detection.

B. Rationale for a Resolution of 30 -100 km

A considerable body of observational evidence indicates the presence of magnetic-related fine structure on scales smaller than the highest resolution yet attained. A few recent examples follow:

  • Magnetic Speckle Imaging: Figure 3 illustrates the presence of Zeeman effect polarization near the diffraction limit of a large ground-based solar telescope (0.2arcsec).

Figure 3: Speckle imaging in combination with polarimetry provides direct evidence for magnetic flux concentrations at or below the angular resolution limit of this large solar telescope: 0.2 arcsec. (Figure courtesy of C. Keller and B.Wilton.)

  • Phase Diversity Reconstruction: With post-facto image reconstruction one may correct for atmospheric seeing. Figure4 reveals many bright features believed to indicate the presence of intense photospheric flux tubes in the intergranular lanes. Some are at the diffraction limit of the telescope (0.16 arcsec, 120km on the Sun). One objective of Sunrise will be to determine the extent to which the intensity in the G-band actually traces the magnetic field.

Figure 4: Phase diversity imaging aided by adaptive optics reveals tiny, bright structures that are unresolved at the diffraction limit of the telescope (0.16 arcsec). These features are associated with concentrations of intense magnetic flux. (Figure courtesy of C. Keller.)

  • Stokes Profile Asymmetry: Departures from symmetry in Stokes Q, U and anti-symmetry in Stokes V spectra signal the presence of gradients of the line-of-sight velocity and magnetic field. Asymmetries are most pronounced where the average flux density is low (Solanki 1993, Sanchez Almeida et al. 1996, Sigwarth et al. 1999), i.e. where convective motions have a larger influence on the magnetic field. Line-of-sight gradients imply dynamical structure on the scale of a photon mean-free-path, which is comparable to a scale height. Observations of extreme asymmetries may be reproduced by atmospheric models having micro-structure on scales small compared to a photon mean-free-path (MISMAs, Sanchez Almeida et al. 1996, 2000). Sunrise is expected to provide much tighter constraints on the distribution of magnetic fields and motions leading to observed asymmetries.
  • Sunspot Fine Structure: Numerous observational studies indicate that much of the fine structure of sunspots remains unresolved, even with the best image reconstruction techniques available today. Fine structure appears to be the essence of sunspot penumbrae, where light and dark filaments are intimately connected with the Evershed flow. In sunspot umbrae, the tiny, bright, transient umbral dots are integral to the energy transport. Sunspot fine structure holds the key to convective transport of energy in a large-scale, strong magnetic field.

C. Interplay of Theory and High Resolution Observations: Solar physics relies increasingly on numerical models to illuminate the physics of varied solar phenomena. For magnetic fields in the solar atmosphere, these models incorporate progressively more sophisticated physical descriptions, yet they still rely on significant assumptions in terms of dissipative processes, boundary conditions, etc. Because of the small-scale nature of the physical processes at work, modeling and simulations of solar magnetic structures, ranging from the weak internetwork flux through sunspot umbrae, all require guidance from observations at the highest spatial resolution possible. Theory, modeling, and numerical simulations also strengthen the case for very small structures:

  • Discontinuous Field Evolution: Driven by convective flows, the footpoint motions of photospheric flux will lead to the development of discontinuities of the magnetic field in the atmosphere above (Parker 1994). The width of such discontinuities in the non-ideal MHD case depends on the conductivity of the plasma. It should be extremely small in the photosphere where diffusion controls the resistivity (Parker 1979). Indeed, MHD numerical simulations in 3D generate discontinuities (Galsgaard & Nordlund 1996).
  • Non-Linear Dynamo Models: Numerical models of magneto-convection generate structure in the magnetic field that is much smaller than the structure in the flows. The dynamo action of convective flows may then by itself create very small-scale structure in the photospheric magnetic field.
  • Diffusive Dissipation in the Photosphere: If molecular diffusion limits the dissipation of magnetic fields in the photosphere, then one expects structure of the field on the scale of 10 km (Schüssler, 1986).

It has often been argued that visible photospheric structure will be limited by photon diffusion to an optical mean-free-path, i.e. about one scale height. This would be true if scattering (i.e. non-LTE effects) entirely dominate the emission and absorption, but collision rates are large in the photosphere, so the emission in most photospheric diagnostic spectral lines is thermal. The result is that structure on scales much smaller than the photon mean-free-path may be detected, even in an optically-thick situation (Bruls & von der Luhe, 2001) Indeed, synthesis of line profiles from dynamical flux tube simulations suggest that visible structure in Stokes V will exist down to size scales of 5 km. In chromospheric lines, non-LTE effects (scattering) often dominate the excitation. Even in this case it is possible to detect structures much smaller than a photon mean-free-path. Under the condition that the structures are optically-thin, tiny, hot, dense structures (i.e., the site of reconnection) would still appear highly localized, albeit with a scattering “bloom” surrounding the dense structure. These considerations, along with the observational and theoretical indications above, strengthen the case for solar observations of the highest angular resolution possible.

D. UV Diagnostics

At balloon float altitude, the Earth’s atmosphere is transparent at some wavelength bands between 200 and 300nm. Figure 5 shows the atmospheric transmission and UV solar irradiance in the range 185-295 nm at a height of 40 km and a latitude of S77.9 on Dec.21 (typical of an Antarctic LDB mission.)

Figure 5: Atmospheric UV transmission (upper panel) and solar irradiance (lower panel) as a function of wavelength and solar zenith angle (sza). Contours and color scale are logarithmic. Lighter shades denote higher transmission or higher flux, respectively. Figure courtesy of D. Marsh, HAO/NCAR.

These images show that, even on an Antarctic LDB mission, a balloon-borne solar telescope will see adequate solar flux for imaging and spectroscopy at 205 nm and 280 nm. The test flight should encounter even higher atmosphere transmission at these wavelengths. Of interest are the highly transmissive wavelengths around the Al I ionization edge at 207 nm and in the vicinity of the Mg II k (279.6 nm) and Mg I (285.2 nm) resonance lines.

The wavelengths around the Al I edge probe the middle photosphere (h  200 km, Vernazza et al. 1981). The opacity there has a strong contribution from spectrum lines. Broadband observations at this wavelength will be very sensitive to the thermal structure at heights where the magnetic diagnostic lines in the visible (i.e. the Fe I lines at 630 nm) form.

In several regards the Mg II resonance lines are better suited as chromospheric diagnostics than the frequently used Ca II resonance lines near 400 nm:

  • The shorter wavelength implies greater sensitivity to temperature.
  • The Mg II resonance lines form higher in the chromosphere than their Ca II counterparts.
  • Because these lines do not share upper levels that branch to other strong transitions (as in the infrared triplet lines of Ca II), the line transfer in the Mg II line wings is highly coherent. Thus, their near wings are suppressed and their chromospheric emission cores are prominent. Mg II k-line spectra thus provide a clear indicator of chromospheric structure and dynamics, and the profile plus its wings sample the photosphere to the upper chromosphere.

High resolution space instrumentation has largely shunned the chromosphere: the important link between the photosphere and corona. Observations of this region are essential to explore the magnetic underpinnings of coronal activity. From the photosphere to the upper chromosphere the pressure falls by a factor of 105-106. Spatially isolated, small flux tubes thus undergo tremendous expansion through the chromosphere. MHD waves may undergo a fundamental change in character while propagating upward through the chromosphere. By providing ultra high resolution, quantitative spectra of the Mg II k-line, Sunrise will fill a very important observational void in our study of the heating and dynamics of the upper solar atmosphere. The freedom from atmospheric seeing will enable:

  • searches for high frequency waves propagating upward from individual flux tubes,
  • recording the dynamics of sunspot fine structure, including umbral flashes,
  • study of the apparently non- or weakly-magnetic oscillations in the internetwork regions of very quiet Sun, including effects of horizontal propagation.
  1. Relationship to Other Programs

With the aid of phase diversity image reconstruction, Sunrise will image the photosphere at 200 nm and the chromosphere in the Mg II k resonance line at 280 nm, and will therefore achieve a spatial resolution of 30 and 40 km on the Sun, respectively. Sunrise will be able to resolve solar structures a factor of 4 - 6 smaller in linear dimension than Solar-B, thus both opening a new realm of solar phenomena to quantitative investigation, and fully resolving (i.e. critically sampling) structures the size of a photospheric scale height. Its capability to perform precision Stokes I spectroscopy in the UV Mg II resonance lines provides a diagnostic capability of the chromosphere unavailable to new large-aperture ground-based telescopes. With its unique capability, the Sunrise program augments worldwide efforts to better understand the small-scale origins of solar variability and the physics of magnetized plasmas in stellar atmospheres.

  1. Technical Approach
  1. Overview

Advantages of a balloon mission over ground-based observations are freedom from atmospheric seeing over a large field-of-view and access to the ultraviolet. Additionally, a two-week LDB mission from Antarctica would provide continuous observing of solar phenomena during one complete disk passage.

The Sunrise instrumentation consists of a main telescope with an aperture of 1m feeding three focal-plane instruments:

  • The Spectrograph Polarimeter (SP): for measurements of all four Stokes spectral profiles, and precision Stokes I spectral measurements of ultraviolet chromospheric line profiles,
  • The Imaging Magnetograph (IMaX): for high resolution, two-dimensional imaging of the photospheric vector magnetic field, and
  • The Filtergraph (FG): for diffraction-limited imaging in selected ultraviolet and visible light wavelength bands

A correlation tracker controlling a tip-tilt steering mirror provides image stabilization and high-precision guiding. An auxiliary full-disk telescope (FDT) provides both a fine pointing signal and full-disk images for interactive target selection.

The German contribution, through the KIS and MPAe institutions will provide:

  • the telescope; including its structure, secondary mirrors, field stop, heat dump, and active wavefront control,
  • the FDT, its CCD detector, and data system,
  • the SP, exclusive of its CCD detectors and DPUs,
  • the FG instrument and data system,
  • the correlation tracker (CT), wavefront control system, and associated control electronics,
  • the optics to distribute the light from the 1-m telescope to its three focal plane instruments, and
  • the instrument control and data storage systems.

The Spanish (IAC) will provide the IMaX.

NASA will provide the lightweight CSiC primary mirror through prior contracts with LMSAL. Under this proposal, NASA would provide:

  • the CCD detectors and Data Processing Units (DPUs) for the SP (HAO), and
  • the balloon gondola and its support systems (power, telemetry, rough and intermediate pointing) (HAO/NCAR)
  1. 1-m Telescope

The clear aperture of the telescope is 1m and the parabolic primary mirror (M1) has a focal length of 2.5m. The elliptic Gregorian secondary mirror (M2) increases the effective focal length of the system to 25m (f/25). Figure 6 shows the optical schematic of the main telescope.

Figure 6: Optical layout of the Sunrise telescope.

A reflective cone of solid copper located at the prime focus rejects about 99% of the incoming flux, but admits a circular field of view of 3.4arcmin (10W of solar radiation) through a 3mm central hole. The flat mirrors M3 and M4 fold the beam so that it is parallel to the optical axis of the telescope and feeds the focal-plane package. The tip-tilt steering mirror M4 is controlled by both the correlation tracker unit and the guider in the FDT. It provides precise pointing and guiding to 0.005arcsec. Baffles and multi-layer insulation (MLI) covering the telescope structure minimize stray light.

Alignment of the primary and secondary mirrors is maintained by a Serrurier truss. The secondary M2 alignment is adjustable in 2 degrees of freedom, and is adjusted dynamically with error signals from the correlation tracker/ wavefront sensor.

Thermal control of the telescope is entirely passive, and the thermal design is greatly aided by the high thermal conductivity of the silicon carbide ceramic (C/SiC) mirrors The primary mirror (M1) radiates the absorbed energy from its backside. The field stop at the prime focus is also passively cooled: heat absorbed by the field stop is removed by conduction to a radiator mounted on the instrument structure.