SUNRISE II: A second look at the Sun

Mission Description    Instrumentation    References

During its two flights in 2009 and 2013, the balloon-borne solar observatory SUNRISE experienced a unique view of our Sun: from a height of more than 35 kilometers and equipped with the largest solar telescope that had ever left Earth, SUNRISE was able to resolve structures with a size of 50 km in the Sun's ultraviolet (UV) light. The journal Astrophysical Journal Supplement now devotes a total of 13 articles to the results of the second flight of SUNRISE. 1) 2)


Figure 1: The solar observatory SUNRISE II is borne by a helium balloon to a float height of more than 35 km (image credit: MPS)

These are complemented by four articles based on data from the first flight that have now been analyzed. In this way, the special edition paints the most comprehensive and detailed picture of the boundary layer between the visible surface of the Sun and its atmosphere in ultraviolet light. The Special Issue reports, among other things, on hot explosions, oscillating fibril-like structures, and the origins of huge plasma flows. The MPS (Max Planck Institute for Solar System Research) in Germany, head of the SUNRISE project, has a key stake in all 17 publications.

Many of the Sun's secrets are revealed only in the UV light that our star emits into space. However, since the Earth's atmosphere filters out most of this radiation, an observing position above this air layer is ideal for solar researchers. The balloon-borne solar observatory SUNRISE offers access to this position - without the immense costs of a space mission. Carried by a huge helium balloon, SUNRISE reaches an altitude of more than 35 kilometers, leaving most of the Earth's atmosphere underneath.

Twice already this concept has proven successful. While SUNRISE witnessed an unexpectedly long activity minimum during its first flight in 2009, in 2013 our star presented itself from a more vigorous side: for almost six days, SUNRISE had an excellent view of sunspots and active regions. MPS researchers published first results from this flight a few months later. More clearly than ever before, the UV data reveal fine structures in the Sun's lower atmosphere only a few kilometers in size such as bright points and long-drawn fibrils near the sunspots.

Since approximately one year, most of the SUNRISE II data has been fully reduced and is now the basis of 13 of the articles published today. In these, the researchers for example elaborate their analysis of the fibril-like structures and determine their shape and lifetime (Figure 2). One of the results: their intensity and width fluctuate on time scales of a few seconds. Such detailed studies were made possible by the high resolution of SUNRISE and the long series of observations.

"With a spatial resolution of 50 to 100 km, SUNRISE provides more accurate observational data in ultraviolet light than any other balloon-borne ore spaceborne solar telescope," says Sami K. Solanki, director at the MPS and head of the SUNRISE mission. In addition, with its two instruments SuFI (SUNRISE Filter Imager) and IMaX (Imaging Magentograph Experiment), SUNRISE looks at a key region of solar research. In the area between the visible surface of the Sun, the photosphere, and the corona, the upper layer of the Sun's atmosphere, researchers hope to find answers to some of the most important open questions of solar physics: how is it possible that with approximately one million degrees the corona is significantly hotter than the photosphere with only 5000 degrees? In which way is the necessary energy from the photosphere transported into the corona and transformed into heat? What is the role of the Sun's dynamic, highly complex magnetic fields? "Everything points to the fact that small-scale and short-lived processes are decisive," says SUNRISE project scientist Tino Riethmüller of MPS.

Discovering these is the mission of SUNRISE. On the first day of the second flight, for example, the observatory witnessed an Ellermann bomb, an explosive but localized increase in radiation intensity and temperature (Figure 3). This phenomenon generally occurs in developing active regions and is regarded as a sign of dramatic reconstruction in the Sun's magnetic field. Magnetic energy is thereby converted into heat, among other things. The simulations complementing the observational data suggest that these changes in the magnetic field architecture originate in the photosphere about 200 km above the visible surface of the Sun.


Figure 2: Left: The Sun's visible surface shows a pattern of so-called granules. They are evidence of hot plasma flows from the Sun's interior, that rise upward, are cooled off and sink down again. Right: In the ultraviolet light from this region long fibril-like structures can be seen (image credit: MPS)


Figure 3: Right: A look at the footprints of coronal loops. Images obtained by NASA's SDO (Solar Dynamics Observatory) on June 12th, 2013 show distinct plasma flows in the Sun's corona. Left: SUNRISE II data documents the magnetic fields that were present on the Sun at the same time and in the same place. Small regions, in which the magnetic polarity is opposite to that of the overall environment prove to be the origins of the loops (image credit: MPS/NASA SDO)

Another process that connects the relatively cool photosphere with the hot corona are coronal loops, impressive arc-shaped plasma flows in the solar atmosphere. Some of them measure up to 100,000 km in size. The starting points of these structures are often found in the vicinity of active regions. The SUNRISE data now allow a precise view of these "footprints". They prove to be places of strong magnetic contrasts: small regions in which the magnetic polarity is opposed to their predominant environment. The interaction of these areas drives mass and energy transport into the atmosphere.

"The data of the two SUNRISE flights are a true treasure trove for solar physics", says Sami Solanki. The analysis of the data will continue for years. In addition, the MPS is currently planning a third flight of the balloon-borne observatory.

SUNRISE is a joint mission led by the MPS (Max Planck Institute for Solar System Research) in Katlenburg-Lindau, Germany. Further partners are the HAO (High Altitude Observatory, Boulder, Colorado), the KIS (Kiepenheuer Institute for Solar Physics,Freiburg, Germany), the Spanish IMaX consortium consisting of the IAC (Instituto de Astrofisica de Canarias), the IAA (Instituto de Astrofísica de Andalucia), the INTA (Instituto Nacional de Técnica Aeroespacial),and the GACE (Grupo de Astronomía y Ciencias del Espacio), the LMSAL (Lockheed-Martin Solar and Astrophysics Laboratory,Palo Alto, California), and NASA's CSBF (Columbia Scientific Ballooning Facility), located in Palestine, Texas.



SUNRISE II balloon-borne solar observatory

The SUNRISE II solar observatory, consisting of a 1 m aperture telescope that provides a stabilized imagery to a UV filter imager and an imaging vector polarimeter, carried out its second science flight in June 2013. The flight provided observations of parts of active regions at high spatial resolution, including the first high-resolution images in the Mg II k line. The obtained data are of very high quality, with the best UV images reaching the diffraction limit of the telescope at 3000 Å after Multi-Frame Blind Deconvolution reconstruction accounting for phase-diversity information. 3)

In spite of the significant advances made by ground-based observations, it remains challenging to accurately take into account or fully remove the effects of atmospheric seeing from observational data. In addition, high-resolution studies from the ground are generally limited to wavelengths with high photon flux and usually to short periods of stable seeing. Thus, in spite of having an aperture smaller than that of the largest ground-based telescopes, the spaceborne SOT (Solar Optical Telescope) onboard Hinode (Tsuneta et al. 2008) has resulted in many advances. These include the discovery of penumbral microjets (Katsukawa et al. 2007), waves carrying copious amounts of energy along spicules (De Pontieu et al. 2007), ubiquitous linear-polarization signals in the quiet Sun (Orozco Suárez et al. 2007; Lites et al. 2008; Lagg et al. 2016), or very fast downflows at the ends of particular penumbral filaments (van Noort et al. 2013). Hence, an even larger solar telescope located above the bulk of the Earth's atmosphere has an extensive discovery space.

The largest solar telescope (along with the Soviet Stratospheric Solar Observatory (Krat et al.1974) 4) to have reached the near-space conditions of the stratosphere is SUNRISE, which had a very successful first flight in June 2009 (for an overview of earlier balloon-borne solar telescopes and results. The data obtained during that first flight of SUNRISE have led to the following discoveries and insights, among many other results (Solanki et al. 2010). 5)

1) First ever spatially resolved images of small-scale intense magnetic flux concentrations in the quiet Sun show that semi-empirical flux tube models provide a reasonable description of such structures.

2) First ever brightness measurements of flux concentrations in the UV at 312, 300 and 214 nm reveal very high intensities, up to a factor of 5 above the mean quiet-Sun brightness at 214 nm.

3) First ever measurements of the rms intensity contrast of granulation in the UV show high values of up to 30%, consistent with numerical simulations. These values provide a direct measure of the efficiency of convective energy transport by granulation.

4) The most sensitive high-resolution time sequences of maps of the vector magnetic field ever obtained reveal abundant, short-lived and highly dynamic small-scale horizontal fields.

5) Ubiquitous small-scale whirl flows are found which drag small-scale magnetic field structures into their centers.

6) Magnetic field extrapolations from SUNRISE/IMaX data indicate that most magnetic loops in the quiet Sun remain within the photosphere. Only a small fraction reaches the chromosphere. Most of these higher-lying loops are anchored (at least in one foot point) in the strong-field elements of the network.

7) Discovery of large amplitude acoustic waves in the quiet solar atmosphere. Such waves were missed in the past, since they are spatially strongly localized and their photospheric sources move significantly within a short time.

8) First detection of horizontally oriented vortex tubes in solar convective features. Such vortex tubes were found to be rather common in solar granules.

9) Discovery that the internetwork magnetic elements continuously move back and forth between a state of weak and strong magnetic field.

10) First determination of the inclination of magnetic elements directly from their position in images sampling different heights. The results reveal that the magnetic elements are nearly vertical, in contrast to inversion techniques that suggested that they were close to horizontal (but in this case were strongly affected by noise).

11) First detection of localized, strongly wavelength-shifted polarization signals in the quiet Sun that are interpreted as supersonic upflows caused by magnetic reconnection of emerging small-scale loops with pre-existing fields.

12) Discovery that 85% of the internetwork magnetic fields stronger than 100 G are concentrated in mesogranular lanes, although there is no particular mesogranular scale.

The data taken during the first flight of the SUNRISE observatory were limited to the quiet Sun, as the Sun was exceedingly quiet for the whole duration of the flight. To be able to probe an active region with the unique capabilities of the SUNRISE observatory, a reflight of the largely (but not completely) unchanged instrumentation was carried out in June 2013.



Instrumentation update: Gondola, Telemetry, Telescope and PFI

Gondola. Structural protective elements on the gondola, such as the crush pad assembly below the core gondola frame or the front and rear roll cages, needed to be replaced, as they were designed to deform and thus to take most of the impact energy to protect the rest of the payload.

The average power consumption during the first flight was much lower than previously estimated, so that the number of solar cells could be reduced to only six panels with 80 SunPower A-300 cells each (from originally 10 panels in 2009). The panels were mechanically re-arranged, reducing the overall width of the instrument. This modification allowed a pre-installation of the solar panels in the integration hall before roll-out, saving precious time on the day of the launch.

The mounting of the electronics racks, carrying all the instrument computers, pointing system computer and power distribution units, was modified for the 2013 flight. Instead of having a 20° tilt toward cold sky as in the previous configuration, the racks were now mounted vertically and were directly attached to the gondola side trusses, providing higher overall stiffness and reduced mass. The reduction in thermal efficiency of this configuration was seen to be acceptable, as results from the 2009 flight indicated that the size of the radiating surfaces are sufficiently large to dump all dissipated heat at moderate temperatures even in this slightly less efficient configuration.

Telemetry. The 2009 flight operations and commissioning were hampered by the early loss of the E-Link high-speed telemetry provided by ESRANGE, which ceased operation after only a few hours when switching from one ground station to another. For the rest of the 2009 mission commanding and health status information was transmitted through a 6 kbit/s TDRSS (Tracking and Data Relay Satellite System) link of NASA. This was found to be clearly insufficient.

In preparation for the 2013 science flight the SUNRISE team helped to qualify the Iridium Pilot/OpenPort system for balloon flights. A dedicated test setup was flown on September 28 and October 7, 2011 from Ft. Sumner, New Mexico, on balloons of the NASA CSBF (Columbia Scientific Ballooning Facility, TX) for several hours each, to verify system integrity during launch, ascent, and operational conditions. The bandwidth transmitted to ground was more than ten times higher than with the TDRSS Omni link. SUNRISE II was one of the first scientific missions to be supported with this new telemetry system. The antenna system on top of the gondola was re-arranged to provide the required 1.4 m clearance to the dome-shaped Iridium transmitter. SUNRISE II reached an average data rate of 100 kbit/s, and the downlink of science data continued even after the payload landed, although the antenna was partly buried in Canadian soil.

Telescope and structure. The telescope and instrumentation needed a thorough cleaning and refurbishment. All mechanisms were disassembled, inspected, cleaned, lubricated, re-assembled, and requalified. The motor and gears of the telescope aperture mechanism were replaced and the mechanism was improved to provide a more reliable detection of the open or closed position. Key elements of the carbon-fiber based structure were inspected, strength tested, and found to be still fully intact. The thermal subsystem, consisting of the heat rejection wedge and its radiators, needed to be refurbished. New second surface mirrors were attached. All telescope mirror coatings were stripped, the mirrors cleaned, their optical quality interferometrically verified and recoated. The secondary mirror needed to be replaced completely, as the Zerodur mirror substrate was damaged during the landing and recovery after the 2009 flight. The mirror was refabricated by the Lytkarino Optical Glass Factory in Moscow, with a performance identical to the original one.

The telescope alignment proved to be a challenge, because a sufficiently large reference flat mirror provided by SAGEM for the 2009 flight was not available this time. Since on the launch site a misalignment of the telescope was detected (the telescope had been tested before departure to Kiruna without a 1 m reference flat, making the performance assessment ambiguous), a readjustment of the M2 position was necessary. Figure 1 is a photo of the SUNRISE II payload (telescope with instruments and electronics mounted in the gondola) hanging from the launch crane at the ESRANGE balloon launch facility.

PFI (Post-Focus Instrumentation). The instrumentation mounted within the PFI platform survived the landing and recovery in good shape. Some cleaning was necessary, but almost all elements could be reused. The PFI was refurbished, but remained technically identical to the one used during the 2009 flight. The scientific instruments were also refurbished and subsequently integrated and realigned.

ISLiD (Image Stabilization and Light Distribution). The ISLiD system (Ref. 9) was not modified relative to the first flight. In order to optimize the end-to-end optical performance of the combined ISLiD–IMaX path, an additional plane-parallel plate was introduced into the converging IMaX feed path in front of the IMaX interface focus. By choosing the inclination and orientation of this plate, residual astigmatism in the ISLiD/IMaX path could be minimized. The plate was coated with high-efficiency antireflective layers on both sides, which were tuned to be unpolarizing at the selected angle of incidence.

CWS (Correlating Wavefront Sensor). The main improvement to the CWS (Berkefeld et al. 2011) for the 2013 flight consisted of a second operating mode. By reading out only the two sub-apertures in the center row (marked red in Figure 4), the bandwidth could be substantially increased and the residual image jitter decreased. Table 1 compares the two modes.

Due to the better performance, the two-subaperture mode was used exclusively during the 2013 flight. Further software improvements included better data logging capabilities and much faster focusing after closing the CWS control loop.


Figure 4: Snapshot of the image recorded by the CWS (Correlating Wavefront Sensor), showing the six sub-apertures displaying the same solar scene. The red frame bounds the two images employed in the two-subaperture mode (image credit: MPS)


6 subaperture mode

2 subaperture mode

Control loop frequency (Hz)
Bandwidth (0db) (Hz)
Measurement accuracy (mas)
Detectable Zernike modes
Residual image jitter (mas rms)

tilt, focus, coma
50 (in 2009)
35 (in 2013)

tilt, focus

25 (in 2013)

Table 1: CWS parameters for the six-subaperture and two-subaperture modes

SuFI (SUNRISE Filter Imager). The SuFI instrument (Ref. 7) was changed very little compared to the SUNRISE I flight. Thus, the mechanical shutter, which had seen more than 150,000 releases during the 2009 flight, was replaced. For the SuFI CCD camera a new power supply unit with very low output noise was developed.

This led to an improvement in SNR (Signal-to-Noise Ratio) performance by a factor of more than ten. In addition, three of the wavelength filters were exchanged. The filter set for observations in the 2140 Å region was replaced by a combination of two different filters, both of which had a FWHM of 210 Å, but had different side band characteristics, which ensured sufficient blocking of longer wavelength radiation. The 3120 Å channel of SUNRISE I was replaced by a combination of a 2795 Å filter with 4.8 Å width and two blockers of 300 Å and 110 Å width, respectively. This combination of filters, whose profile is plotted in Figure 5 of Riethmüller et al. (2013a), is centered on the Mg II k line and gets minimal contribution from the Mg II h line located in the filter profile's wings. Observations in the Ca II H line were possible through two different filters centered at 3968 Å. In addition to the 1.8 Å wide filter already used in SUNRISE I, a 1.1 Å wide filter was available, which provided better isolation of the contributions from higher atmospheric layers. In the following, these channels are referred to as the 3968w ("w" for wide wavelength band) and 3968n ("n" for narrow band) channels. The 3880 Å channel, which was used in SUNRISE I, was sacrificed for this purpose.


Figure 5: Observed Mg II h&k spectrum in plage (solid gray) and the same multiplied with the SuFI filter transmission curve (solid black; multiplied by 10 for better display). The SuFI Mg filter transmission profile is overplotted with the corresponding y axis on the right (dashed red), image credit: MPS

IMaX (Imaging Magnetograph eXperiment). The IMaX instrument (Ref. 8) on the 2013 flight was also very similar to the version flown on SUNRISE I, although a number of smaller changes and updates had been made. The parts replaced included the FPGA (Field Programmable Gate Array) in the proximity electronics, the LiNbO3 etalon, one of the CCD cameras, the collimator doublet, the LCVRs (Liquid Crystal Variable Retarders), the power supply for the LCVR heaters (the heaters on the SUNRISE I flight did not have enough power to bring the instrument to its nominal operational temperature) and the PD (Phase-Diversity) plate (the original broke upon landing after the SUNRISE I flight). Most of these parts were replaced by nearly identical ones. In addition, the windows of both CCD cameras were removed in order to avoid fringes that had been observed during the SUNRISE I flight.



Description of the Mission

SUNRISE II was flown on a zero-pressure stratospheric long-duration balloon, launched and operated by the CSBF (Columbia Scientific Ballooning Facility). It was launched on 2013 June 12 at 05:37:53 UT (07:37:53 local time) from ESRANGE (67.89°N, 21.10°E) near Kiruna in northern Sweden on a cloudy, but perfectly windstill day. It reached a float altitude of 37.1 km after an ascent lasting approximately 3.5 h. It then drifted westward at a mean speed of 35.3 km/hr and at a mean altitude of roughly 36 km. Compared to SUNRISE-I it travelled more to the south and somewhat faster across the Atlantic and Greenland, so that it reached the peninsula of Boothia in northern Canada on 2013 June 17 where the flight was terminated at 11:49:24 UT.

After 122.3 hr at float altitude, the balloon was cut off and the payload descended suspended from a parachute, reaching the ground at 70.08°N, 94.42°W about one hour later. It landed relatively softly but tipped over forward, so that the front ring of the telescope and the radiators of the heat rejection wedge thermal control system were damaged. Otherwise very little serious damage was incurred. Figure 6 (a) shows the flight path of SUNRISE-II overplotted on a map. Also plotted, for comparison purposes, is the flight path taken by SUNRISE-I almost exactly four years earlier. Figure 6 (b) displays the height profile of the flight, with the day–night cycles being visible (although the Sun never quite set on the payload at float altitude). Rapid, although usually not very large, increases in height, for instance on June 15 starting at 05:04 UT, are due to ballast drops. The elevation of the Sun as visible from SUNRISE-II is overplotted. Clearly, the payload remained in direct sunlight during the entire flight, although the Sun was just over the horizon around local midnight.

When at float altitude, the payload was above 99% of the Earth's atmosphere, virtually seeing-free observations were possible all the time. Also, the balloon stayed above most of the ozone in the Earth's atmosphere, allowing high-resolution imaging in the UV at 2140, 2795, and 3000 Å, although the residual atmosphere did require excessively long observing times at 2795 Å and, to a smaller extent, at 2140 Å.

E-Link worked until 2013 June 13 at 1:05 UT, i.e., until well after commissioning was completed. Two events occurred after commissioning that impacted the science operations.

(1) The temperature controller of the IMaX etalon failed on June 12, 2013 at 23:55 UT, although possibly problems in the communications between IMaX subsystems was the real reason for the failure.

(2) On June 13, 2013 at 7:30 UT, about 23 hr after the observations had started, the highly reflective front face of the heat rejection wedge—a glued, 0.1 mm thin second surface mirror—failed. Due to increased absorption the heat rejection wedge temperature rapidly increased each time the telescope was pointed at the Sun, eventually exceeding the temperature sensor measurement range (whose upper limit lay at 130ºC) within a number of minutes (the exact number depended somewhat on the time of day). Therefore, all observing sequences had to be shorter than this interval after this incident, with roughly half an hour in between observing sequences to allow the heat rejection wedge to cool down again. However, the optical quality of the data obtained after the failure of the heat rejection wedge remained the same as before, as a careful inspection revealed.


Figure 6: (a) The flight paths of the 2009 (red curve) and 2013 (blue curve) SUNRISE science flights overplotted on a map of the northern Atlantic. The semi circles mark latitudes of 60º, 70º, and 80º, respectively. (b) The SUNRISE-II float altitude vs. time as recorded by the CSBF-provided SIP (Support Instrumentation Package), from shortly before launch up to the time of cut-off from the balloon (black curve, referring to the left axis). Also plotted is the solar elevation angle, as seen from SUNRISE-II (in red, referring to the axis on the right).


Overview of the Data Recorded during the 2013 Flight of SUNRISE

The total length of time over which observations were made was 122 hours, during which period SuFI recorded 300 GB (60,806 images), while IMaX acquired 68 GB of data (48,129 images). During 16% of the total time at float altitude the CWS loop was closed. The longest time series of SuFI data covers 60 minutes, while the longest currently reduced IMaX time series lasts approximately 17 minutes.

SuFI observed in a variety of modes, differing mainly in the wavelengths sampled. The various modes were recorded at different times of day. Thus the two shorter wavelengths, at 2140 and at 2795Å, were only recorded close to local noon, when the Sun was the highest in the sky and the absorption due to the atmosphere was minimal. A brief summary of the recorded SuFI data is given in Tables 2 and 3. The plate scale per pixel of SuFI depended on the wavelength and was in the range 0'' 01983 – 0'' 02069.

SuFI mode

Fraction of observing time

5l: 2140, 2795, 3000, 3968w, 3968n Å
4l: 2140, 3000, 3968w, 3968n Å
3l: 3000, 3968w, 3968n Å
2l: 3000, 3968w Å
1l: 3968w Å


Table 2: SuFI Observing Modes

Legend to Table 2: 3968w refers to the SuFI channel observed through the 1.8 Å wide filter centered on the line core of Ca IIH, while 3968n refers to the channel observed through the 1.1 Å narrow Ca filter.

Start Time (UT)

Duration (s)

Filter (Å)

Experiment times (ms)

Cadence (s)



IMaX mode

12.06. 23:39
13.06. 01:10
17.06. 01:43
16.06. 04:37
12.06. 19:59
16.06. 00:43
17.06. 06:02
16:06. 05:04
16.06. 00:10
16.06. 05.54
14.06. 02:58
17.06. 02:37


3000, 3968w, 3968n
3000, 3968w, 3968n
2140, 3000, 3968w, 3968n
3000, 3968w, 3968n
3000, 3968w, 3968n
2140, 3000, 3968w, 3968n
3000, 3968w, 3968n
3000, 3968w, 3968n
3000, 3968w
3000, 3968w, 3968n
1 x 3000, 5 x 3968w
3000, 3968w, 3968n

500, 100, 500
500, 100, 500
2500, 50, 100, 500
200, 100, 500
50, 100, 500
2000, 50, 100, 500
50, 100, 500
200, 100, 500
50, 100
200, 100, 500
500, 100
50, 100, 500





Table 3: List of the Longest SuFI Time Series of Active Region Observations

IMaX data were dominantly obtained in the V8-4 mode, meaning that the full Stokes vector of the Fe I 5250.2 Å line was recorded at eight wavelength positions, with four accumulations at each wavelength. The wavelengths sampled were centered on -120, -80, -40, 0, +40, +80, +120, and +227 mÅ from line center. The last of these samples a continuum position between the g= 3 line and its neighboring Fe I line at 5250.6 Å. Each exposure lasted 250 ms, so that the cadence achieved with the IMaX V8-4 mode was 36.5 s. The plate scale per pixel of IMaX remained unchanged at 0'' 05446.

A relatively small amount of IMaX data were also recorded in the L12-6 mode, in which only Stokes I and V are recorded at 12 wavelength points in the line, with six accumluations per wavelength point.



Data Reduction of SUNRISE II (Ref. 3)

Due to the changes in the instrumentation as well as the experience gained from SUNRISE I, the data were reduced using a modified reduction pipeline.

SuFI: The SuFI data, as on the first flight, were acquired in PD (Phase-Diversity) mode, i.e., one half of the sensor was located in the focal plane, while the other half imaged the same part of the solar surface as the first half, but with a fixed offset in the focus direction. This configuration is necessary, since aberrations of the telescope, although small enough to allow for diffraction limited performance at visible wavelengths, are not negligible at the much shorter UV wavelengths. The PD technique in principle allows for the determination and subsequent removal of these aberrations, provided one of the pair of images is in or very near focus according to Gonsalves & Chidlaw (1979) and Paxman et al. (1992).

Therefore, after the traditional darkfield and flatfield corrections, the data were restored using the MFBD (Multi-Frame Blind Deconvolution) wavefront sensing code, which allows for an arbitrary PD term to be present in the data. Due to the highly non-telecentric configuration of the beam, however, the image scale of the defocused image was slightly different from the image scale in the focal plane, so that the influence of the tip-tilt components of the wavefront aberrations in the defocused channel differed from those in the focal plane. This effect was accommodated by restoring the frames in the "calibration mode" of the MFBD code, where all fitted wavefront modes are constrained to be the same for the two channels, with the exception of the tip-tilt and focus terms, which are allowed to vary, but with the difference between the focus and defocus channels constrained to have the same value for all PD pairs in the data set. Due to the relatively low frame rate of the SuFI camera, only four PD pairs, recorded in approximately twenty seconds, could be combined to restore each frame, beyond which the solar evolution started to degrade the result.

Although this method worked well on some of the data, it did not always restore the data to the same quality. This is believed to be the result of residual pointing errors during the relatively long exposure time, that blurred the image in a way that is not consistent with the model, which assumes the PSF (Point Spread Function) to be produced by a single, static wavefront. The azimuthally averaged Fourier power spectrum of the best frames, however, shows power significantly above the noise floor, at wave numbers up to 85% of the true diffraction limit of Dtelescope /lambda (DL, which is slightly better than the criterion of Dtelescope / (1.22 λ) proposed by Rayleigh (1879) and generally adopted for considering individual features to be resolved.

IMAX: The reduction applied to the data from the 2009 flight was explained in detail by Martínez Pillet et al. (Ref. 8). Consequently, here the focus is on changes in the data reduction with respect to the reduction of the first flight's data, necessitated by the replacement of several components of IMaX. These replacements influenced various instrumental effects of IMaX such as the occurrence of fringes, ghost images, and the instrument's stray-light behavior.

After the data were corrected for dark current and flat-field effects, residual excess power due to interference fringes was filtered out in the Fourier domain. Before and after each observing run during the flight, a PD measurement was obtained by inserting a glass plate into the optical path of the first camera in order to record pairs of defocused and focused images, which then allowed retrieval of the system's PSF. The observational data were reconstructed by applying a modified Wiener filter constructed from the PD PSF. By spatially replicating the images before performing operations in the Fourier domain, a reduction of the FOVs as a side effect of a necessary apodization could be avoided. Hence the effective FOV of the IMaX data from SUNRISE II is 51'' x 51'' (or 51 arcsec x 51 arcsec).

The instrumental polarization was removed from the observations by a demodulation matrix determined from the pre-flight polarimetric calibration. In contrast to the first flight, this time the FOV dependence of that matrix was taken into account. Since the on-ground polarimetric calibration did not include the main mirror of the telescope and because the thermal environment during calibration was different to the inflight situation, cross-talk with Stokes I had to be removed from the data. Then the images of the two cameras were co-aligned and merged.

Finally, we determined the spatial mean Stokes I profile and subtracted 25% of the mean profile from the individual profiles to correct the data for 25% global stray light. This is equivalent to deconvolving the IMaX data with a constant stray light not varying over the FOV.

The physical quantities of the solar atmosphere were then retrieved from the reconstructed Stokes images via the SPINOR (Stokes Profiles INversion O Routines), Frutiger et al. 2000), which uses the STOPRO routines to solve the Unno–Rachkovsky equations. A relatively simple inversion strategy was applied to get robust results: three optical depth nodes for the temperature (at logτ = -2.5, -0.9, 0) and a height-independent magnetic field vector, line of sight velocity, and micro-turbulence. The synthetic spectra were not only calculated for the Fe I 5250.2 Å line but also for its neighboring lines (Co I line at 5250.0 Å and the Fe I line at 5250.6 Å) to include their influence. The spectral resolution of IMaX was taken into account by convolving the synthesized spectra with the spectral transmission profile of IMaX measured in the laboratory during a pre-flight calibration campaign.

The SPINOR inversion code was run for ten iterations, then the output maps of the individual parameters were spatially smoothed and used as the initial guess parameters to a further run of ten iterations. The intermediate smoothing of the free parameters of the inversion was introduced to get rid of spatial discontinuities in the physical quantities that can occur if the inversion gets stuck in local minima of the merit function at individual pixels or groups of pixels. This procedure was repeated five times (implying 50 iterations in all), with a gradually decreasing amount of smoothing. Finally, the resulting LOS velocity maps were corrected for the etalon blueshift, which is caused by the collimated setup (Ref. 8).


Figure 7: HMI (Helioseismic and Magnetic Imager) continuum map (left image) and HMI magnetogram (right image) recorded on 2013 June 12 at 23:39 UT, at the beginning of the SUNRISE II time series. The x and y scales are in helio-projective cartesian coordinates and are in arcsec, with the origin located at the solar disc center. The outer blue box indicates roughly the FOV of the IMaX instrument, while the inner blue box does the same for the SuFI instrument. The animation of this figure shows the evolution of the active region for three days around the SUNRISE II observations (image credit: SUNRISE II Team)


Examples of Data and Science Results

NOAA AR 11768: The longest time series obtained by SUNRISE II was focused on NOAA AR 11768 at cos θ = μ = 0.93, where θ is the heliocentric angle. The observations were made on June 12 2013, starting at 23:39 UT, about 1.5 days after the initial appearance of the active region, at a time when magnetic flux was still emerging. At that time the active region had developed a full-fledged leading polarity sunspot, while the following polarity was mainly concentrated in a string of pores.


Figure 8: Images recorded by SuFI on June 12, 2013 at 23:46 UT after multi-frame blind deconvolution incorporating phase diversity. Plotted are (a) the intensity in the filter centered on 3000 Å, (b) intensity in the broad Ca II H channel and, (c) in the narrow Ca II H channel. All intensities are given in units of the mean quiet-Sun intensity. The coordinates are given with respect to the lower left corner of the IMaX images (Figure 10). Note that in this figure the plots at various wavelengths are aligned with each other to an accuracy of only roughly 1'' (image credit: SUNRISE II Team)

The active region is illustrated in Figure 7 at the time of the SUNRISE II observations. The IMaX and SuFI FOVs are overplotted on a Helioseismic and Magnetic Imager (HMI) continuum image and an HMI magnetogram and are indicated by the blue boxes. IMaX covered most of the following magnetic polarity pores, including the largest one. The largest pore had a complex structure and on one side showed a feature with a roughly penumbral brightness exhibiting elongated structures. In the course of the further development, however, this feature did not develop into a proper penumbra (as can be deduced from the animation of Figure 7). The IMaX FOV also contained some of the leading polarity flux, mainly in the form of facular magnetic elements. SUNRISE II also caught a region of emerging flux, to be found in the central southern part of the IMaX FOV and partly in the much smaller SuFI FOV, which captured mainly the region between the two polarities, including one of the emergence events.

The animated Figure 7 shows the evolution of the active region from the very first emergence 1.5 days prior to the start of SUNRISE II observations of this region to approximately 1.5 days after the end of the observations of this region. A significant amount of flux appeared in the day after the SUNRISE II observations, which led to the formation of a proper following polarity sunspot, as well as to an increased size of the leading polarity spot, mainly through the coalescence of pores, many of which already carried a piece of penumbra on one side prior to merging. The blue boxes overlaid on the images in the animated figure are tilted by the rotation angle relative to solar north and the roughly 5° of rotation during the course of the SuFI time series is taken into account. Note that the animated Figure 7 runs at a slower speed during the time of the SUNRISE II observations of this region.

Sample SuFI Data: The SuFI instrument during the second flight of SUNRISE provided diffraction-limited images at 3000 Å and at 3968 Å in both the broader and the narrower channel centered on the core of the Ca II H line. An image at each of these wavelengths is plotted in Figure 8. Note that during this, the longest time series obtained by SuFI, no data in Mg II k, nor in the 2140 Å channel were obtained due to the low solar elevation at that time and the consequent high atmospheric absorption at these wavelengths. Strongly different structures are seen at 3000 Å and in the Ca II H line, with the shorter wavelength displaying granulation (as had already been noticed in the SUNRISE I data, e.g., Hirzberger et al. 2010; Solanki et al. (Ref. 5), but also bright points (Riethmüller et al. 2010), pores with internal fine structure, and a bright elongated granule located at approximately 30''–32'' in the x-direction and 15''–16'' in the y-direction. This last feature is found to be associated with magnetic flux emergence. The two brightest features in the 3000 Å image are located near the two ends of this structure; on its left it brightens to the level of a strong bright point, while there is a bright facular region right at the edge of the FOV in the continuation of the elongated granule on its right side.

The surprising lack of difference between the broader and narrower Ca II H channels in Figure 8 suggests that they were observing very similar layers of the Sun. Possibly, the filter selecting the narrower channel had drifted in wavelength (e.g., due to temperature changes in the instrument) and was not centered exactly on Ca II H3. Due to the similarity between the broad and the narrow Ca II H channels, in the following the discussion will concentrate on the narrower channel.

In the Ca II H images the dominant structures are slender fibrils directed roughly across the SuFI FOV, many of which seem to emanate from the large pore just to the left of the SuFI image. They appear to be similar to those observed earlier by Pietarila et al. (2009). Although most of the filaments seem to roughly follow the same direction, a number of them do cross each other. This is particularly clearly seen close to the location of the emerging magnetic flux mentioned above, at around (30''–34'', 13''–15''). Very likely this is due to fibrils belonging to previously existing magnetic field overlying those associated with the emerging flux that are directed differently. The lengths of these fibrils are hard to determine, due to inhomogeneities that may be present in the background, but some seem to extend over a fair portion of the width of the SuFI images. Strikingly, the fibrils are also constantly evolving and in motion. The lengths, along with other properties of these fibrils, are studied in a further paper in this special issue, while the dynamics of the fibrils, in particular wave modes travelling along them, are investigated by Jafarzadeh and by Gafeira.

The brightenings seen at 3000 Å in connection with the flux emergence are also visible in Ca II H. The magnetic flux emergence caught by IMaX and SuFI is investigated by Centeno and Danilovic. At a later stage during the time series this brightening develops further into a small flare that engulfs the emerging flux region and is prominently visible in the Ca II H line.

SuFI on board SUNRISE II also obtained the first high-resolution images in the Mg II k line, which have already been published in two earlier papers (Riethmüller et al. 2013a; Danilovic et al. 2014). Images taken in the Mg II k filter were found to have considerable similarity with nearly simultaneously recorded Ca II H images, although some distinct differences were also found. These included a 1.4–1.7 times larger intensity contrast and a more smeared appearance. Examples of Mg II k images of the quiet Sun and of a weak plage/enhanced network region are shown in Figure 9. These images were exposed for 50 s to counter the strong atmospheric absorption at this wavelength and to obtain a sufficient SNR. The images are based on SuFI level 3 data, i.e., data that have been PD reconstructed using wave-front errors averaged over the whole time series of images. See Hirzberger et al. (2011) for more details on reduction and PD reconstruction of SuFI data.

It was concluded that some of these differences may be caused by the much longer integration time needed to record the Mg II k images due to the strong ozone absorption at that wavelength, while others are likely intrinsic. Possible reasons for the intrinsic differences given by Danilovic et al. (2014) include greater formation height and greater formation height range of the Mg line (the latter due to the rather broad Mg filter) and the stronger response of the emission peaks of Mg II k to temperature.


Figure 9: Images recorded by SuFI in the Mg II filter. Panel (a) displays a part of the quiet Sun recorded on June 13, 2013 at 12:52:50 UT, while in panel (b) a weak plage is shown that was observed on June 12, 2013 at 12:50:48 UT. Both images were recorded exactly at disc center and correspond to SuFI level 3 data (see the main text for details) and have been PD reconstructed using an averaged point-spread function (image credit: SUNRISE II Team)

Sample IMaX Data and Inversion Results: The continuum intensity in the visible recorded by IMaX at +227 mÅ from the line center is plotted in Figure 10 (a). IMaX images contain a larger number of pores than SuFI due to the bigger FOV. As expected for the continuum intensity at 5250 Å, bright points are much less prominent than at 3000 Å (Figure 8 (a)), although the fine structure in the largest pore is clearly visible, with some of the bright "umbral dots" being among the brighter structures in the image. Also clearly visible in that figure is the long elongated granule already pointed out in the 3000 Å image (located at (30'', 15'') in the IMaX FOV). Another similar granule is found to its right. Both are associated with magnetic flux emergence.

The bright points in the line core image displayed in Figure 10 (b), often forming connected chains filling the intergranular lanes, possess a high contrast. Here, what we refer to as the line core intensity is simply the intensity at the central IMaX wavelength point, which is nominally located at the wavelength of the Fe I 5250.2 Å line core. Chains of such bright points surround most of the pores, but the most prominent bright feature in this image is the strong brightening located between the two elongated granules, i.e., between the two small flux emergence events. The line core of Fe I 5250.2 Å obviously forms below the height of formation of the radiation sampled by either of the two SuFI Ca II H filters, as the iron line core image still shows faint vestiges of granulation and no signs of fibrils. All quantities in Figure 10 are normalized to the respective intensity averaged over two areas within the joint IMaX and SuFI FOV with near-quiet-Sun conditions, i.e., a relatively low magnetic flux density. We call this averaged "quiet-Sun" intensity, IQS (Intensity Quiet Sun). Note that the granule pattern changes to that of reversed granulation if instead of being divided by IQS, as is the case in Figure 10 (b), the line core intensity is divided by the local continuum intensity, i.e., I+227, at each pixel.

Stokes I, Q, U, and V in the red flank of 5250.2 are plotted in Figures 10 (c)–(f). The differences between Figure 10 (c) and Figures 10 (a) and b are due to the formation height of the line flank (at +40 mÅ) lying between those of the other two wavelengths and the fact that the intensity in the line flank, i.e., in Figure 10 (c), is very sensitive to the Doppler shift of the spectral line. The intensity in the line core, as plotted in Figure 10 (b), is also affected by Doppler shifts, but to a much lesser extent than the line flanks.

Maps of the best-fit parameters obtained from the inversion of the IMaX Stokes vectors illustrated in Figure 10, are presented in Figure 11. The panels in the upper row display the temperature at the three nodes at which it was determined, log(τ) = 0, log(τ) = -0.9 and log(τ) = -2.5; the lower panels give the magnetic field strength, B, the inclination of the magnetic field relative to the line of sight, γ, and the line-of-sight velocity, νLOS.

The results of the inversions reveal some conspicuous features. Among these are the strong blue- and redshifts in the granulation (note that we have not removed the p-modes from the inverted data, so the velocities are influenced by such oscillations). The velocity amplitude is larger than in the granulation shown by Solanki et al. (Ref. 5). An important reason for this is that scattered light has been removed from the present data, while the data analyzed by Solanki et al. (Ref. 5) were still affected by scattered light. Some discrepancy can also be introduced because the velocities displayed by Solanki et al. (Ref. 5) were obtained from a Gaussian fit to line profiles sampled at four wavelengths, while the velocities presented here were provided by the SPINOR inversions of line profiles sampled at seven wavelength points. One of the emerging flux patches is found to be associated with a strong blueshift, much stronger than the other emerging flux feature.


Figure 10: IMaX observables recorded on June 12, 2013 at 23:46 UT after PD reconstruction. Plotted are (a) the continuum intensity, I+227, (b) line core intensity (i.e., intensity in the wavelength channel nominally located at the line core), (c)–(f) the Stokes I, Q, U, and V values at +40 mÅ from the line center. All quantities are given in units of the respective quiet-Sun intensity, IQS (image credit: SUNRISE II Team)


Figure 11: Best-fit atmospheric parameters deduced from the inversion of the Stokes vectors recorded by SUNRISE/IMaX and shown in Figure 10. Upper row of panels, (a)–(c): temperature T (log τ = 0), T (log τ = -0.9) and T (log τ = -2.5). Lower row, (d)–(f): magnetic field strength, B, magnetic field inclination, γ, and the line-of- sight velocity, vLOS. Positive velocities indicate downflows. γ is relative to the line of sight direction (image credit: SUNRISE II Team)

Pores: One noticeable feature of the inversion results is that the photospheric field strength in pores often reaches 2500 G (comparable with the highest values found from Hinode data after deconvolving with the PSF; Quintero Noda et al. 2016). In some places higher field strengths are reached, but these cannot be trusted due to the limited wavelength range of the observations and the fact that the continuum point at +227 mÅ starts to be affected.

Another striking feature is that, whereas all pores are clearly cooler than their surroundings at the two lower heights (log t = 0, -0.9), only the larger pores are clearly distinguishable from their surroundings in the temperature map at log τ = -2.5. The smaller pores can hardly be separated from their surroundings. In contrast, the line core image in Figure 10 shows a darkening also at the locations of the smallest pores. However, this is mainly due to the lower continuum intensity at these locations. Also, the line core is affected by the Zeeman effect and to a smaller extent also Doppler shifts, so that it is not an unalloyed measure of the temperature. In order to test whether the inversions are giving a reasonable representation of the temperature stratification in pores we consider the SuFI Ca II H channel.

To compare SuFI and IMaX data, the images at individual SuFI wavelengths were aligned to sub-pixel accuracy with the most similar IMaX images. For 3000 Å this is the 5250.4 Å continuum, while the Ca II H channel is compared with the IMaX line core, normalized to I+227 at each pixel individually (then the images look more similar).

Thumbnails of T at log τ = 0, -0.9 and -2.5 as well as of IMaX continuum intensity, I+227, and the intensity in SuFI 3000 and 3968n Å channels ("n" is for the narrow Ca II H channel are plotted in Figure 12 for two small pores in the top-right part of the SuFI FOV, and in Figure 13 for a single, very small, somewhat weaker pore in the lower part of the SuFI FOV. Overlaid on both pores are contours of the 1400 G level of the magnetic field. In the following, we consider this contour as an independent boundary of the pore, or more exactly of the magnetic concentration underlying the pore.

1400 G is found to be a good compromise, on the one hand, to isolate a pore from the bright points in its neighborhood while, on the other hand, keeping as much of the pore as possible within the contour. For example, if we take 1200 G or 1300 G, then the bright points often found in the immediate vicinity of pores are included in the contours of some of the pores. If, however, we use a larger value, such as 1500 G or 1600 G, then we include less of the dark part of the pore in the contour. To illustrate this we have added also contours of 1200 G and 1600 G besides those of 1400 G in Figure 12. The 1200 G contour includes bright extensions from the smaller (lower left) pore in this image, while the 1600 G contour misses a significant fraction of the upper right pore.

Interestingly, the 1400 G contour often lies outside the actual pores as seen in, e.g., I+227 or the SuFI 3000 Å channel. When comparing with I+227, a part of this difference may have to do with the expansion of the field with height. However, the difference in size is just as large when considering the temperature at log t = -0.9, which is much closer to the height at which the magnetic field is determined. Hence, these images demonstrate that pores are surrounded by regions of strong magnetic field that extend well beyond the visible boundaries of the individual pores (see also Martinez Pillet 1997). Of course, the magnetic field, being determined from multiple recordings, is more susceptible to jitter, and may not have the same high resolution as the individual images it is compared with. This extension of strong fields beyond the boundary of the pore still holds even if we consider a threshold of 1500 G or even 1600 G, even though some of the dark parts of the pores no longer lie within the 1600 G contour.

Figures 12 and 13 confirm that the temperature in the lower two inversion nodes are clearly lower than in the surroundings, but are rather similar in the upper node. Similarly, in I+227 and at 3000 Å the pores are all clearly dark, while in Ca II H the picture is more mixed. Whereas the pore in Figure 13 is clearly as bright as its surroundings and even has a particularly bright intrusion, the pores seen in Figure 12 both appear to be somewhat darker than their immediate surroundings.

We note, however, that the immediate surroundings of the pores in the Ca images are relatively bright (Figure 8). The pores are often surrounded by considerable magnetic flux that produces a local brightening. In radiation coming from the lower photosphere, some of this is visible as bright points, or particularly bright granules or granule walls (similar to faculae near the limb, e.g., Carlsson et al. 2004; Keller et al. 2004; Lites et al. 2004). In the low chromosphere, however, this concentration of magnetic flux produces enhanced brightenings in the form of mainly short fibrils surrounding the pores. The pores are not substantially darker than other, less bright parts of the solar surface in the Ca II H filter (Figure 8).

To study this more quantitatively, we plot the corresponding histograms in Figures 14 and 15. Each set of histograms represents either the three temperatures, at log τ = 0, -0.9, -2.5 (left panels), or the three intensities (IMaX I+227, SuFI 3000 Å, SuFI 3968n Å; right panels) studied here. All intensities and temperatures in Figures 14 and 15 are normalized to patches of comparatively quiet Sun (i.e., with low amounts of magnetic flux) near the top left and the bottom of the SuFI FOV.


Figure 12: Blow-ups of a region containing two pores within the SuFI FOV. Upper row of panels, from left to right: temperature T (log τ = 0), T (log τ = -0.9) and T (log τ = -2.5). Lower row, from left to right: IMaX continuum intensity, I+227, SuFI 3000 Å and SuFI 3968n channel. The solid contours around the pores mark a field strength of 1400 G. For illustration purposes, we have added contours corresponding to 1200 G (dashed) and 1600 G (dotted). In the lower panels the contours have been given different colors to enhance clarity (image credit: SUNRISE II Team)

In all pores the temperatures in the two lower layers and the intensities at the wavelengths formed deeper (IMaX I+227 and SuFI 3000 Å) are predominantly lower than in the quiet Sun. The pixels in which they are higher are often the hot walls of neighboring granules that are seen through a strong magnetic field. The same is true for the intensities in the two deeper-forming wavelengths. The temperature at the top node and the brightness in the Ca II H channel behave quite differently, however. The temperature lies below the quiet-Sun value for only a small fraction of the points. The same is also true for the Ca intensity, for which the ratio of hotter and brighter points to cooler and darker ones is even more extreme.

Hence, although the small pores partly remain visible as dark features relative to their immediate surroundings in the lower chromosphere, they are approximately as bright as, or even brighter than, the quiet Sun at those layers. The most extreme enhancement is displayed by the smallest of the three pores, with the largest being the coolest and darkest (by a small amount) of these three.

The largest pore in the IMaX FOV (and partly in the SuFI FOV) does remain considerably darker than the quiet Sun, even in the Ca images, which unfortunately cover only a small part of the whole pore. Nonetheless, the behavior of the various pores taken together suggests that the difference in the vertical temperature gradient within these pores to that in the quiet Sun depends significantly on the size of the pore. The photospheric temperature gradient is flattest for the smallest pores, with the temperature starting nearly 1000 K below the average quiet-Sun value at the solar surface and reaching the quiet-Sun value in the upper photosphere. For large pores, however, the temperature gradient appears to be more similar to that in the quiet Sun.

Inhomogeneities in brightness and temperature within the pores also decrease with height. Thus, in Figure 12 the umbral dot-like brightenings in the pore at the upper right of the frame are clearly visible in the I+227 at 5250.4 Å, but only very weakly visible at 3000 Å, which is formed only about 50 km higher, and cannot be seen at all at 3968 Å. This suggests that such inhomogeneities are likely of convective origin, as proposed by Parker (1979), Choudhuri (1986), and Schüssler & Vögler (2006) for umbral dots. In sunspot umbrae a similar restriction to low heights of the thermal enhancements/brightenings associated with umbral dots has been found by, e.g., Socas-Navarro et al. (2004) and Riethmüller et al. (2008, 2013b).


Figure 13: The same as Figure 12, but for a single small pore (image credit: SUNRISE II Team)


Figure 14: (a) and (c) Histograms of T (log τ = 0), T (log τ = -0.9) and T (log τ = -2.5 within the 1400 G contours of the two pores displayed in Figure 12. (b) and (d) Corresponding histograms of IMaX I+227, SuFI 3000 Å, SuFI 3968n Å. Panels (a) and (b) represent the larger pore, in the upper right part of the images in Figure 12, panels (c) and (d) the pore in the lower left. Histogram colors distinguish between the different log τ layers, or wavelengths, as labelled in the individual panels (image credit: SUNRISE II Team)


Figure 15: The same as Figure 14, but for the small pore displayed in Figure 13 (image credit: SUNRISE II Team)

In summary, the second science flight of SUNRISE in 2013 June allowed this balloon-borne solar observatory to obtain the first seeing-free observations of an active region close to the diffraction limit of the 1 m diameter telescope. These data are rich in information about a variety of solar phenomena at very high spatial resolution and at wavelengths that partly cannot be accessed from the ground.



SUNRISE I Observatory Instrumentation (Telescope, SuFI, IMaX, ISLiD, CWS)

The SUNRISE PFI (Post Focus Instrumentation) consists of four units, two of which are science instruments, the other two are system units for image stabilization and light distribution.

Note: The description of the SUNRISE I payload has been added for a better instrument reference.


The telescope is a Gregory-type reflector with 1 m clear aperture and an effective focal length of close to 25 m. A heat-rejection wedge at the prime focus reflects 99% of the light from the solar disk off to the side, reducing the heat load on the post-focus instruments to approximately 10 W. The secondary mirror is actively controlled in three degrees of freedom to compensate for thermoelastic deformations of the telescope during flight. The post-focus instrumentation rests on top of the telescope (Figure 18).

The telescope uses innovative mirror technologies, active in-flight alignment and image stabilization systems. Thanks to its high spatial resolution of less than 0.1 arcsec, the telescope is capable of resolving structures smaller than 100 km on the solar surface. 6)


Figure 16: Optical design of the SUNRISE telescope (image credit: MPS)

The SUNRISE telescope Gregory-type reflector consists of a parabolic primary mirror (M1) with 2.5 meters focal length and an elliptical secondary mirror (M2). These mirrors are built into a low-mass and very stable carbon fiber structure. Two plane folding mirrors (M3 and M4) behind the primary mirror redirect the light from the telescope to the PFI (Post Focus Instrumentation) sitting piggy-back on the telescope.

Solar telescopes require special measures to protect their instruments from the Sun's radiation. During the observations, nearly 1 kW solar radiation is concentrated on a disk of about 22 mm diameter in the first focal spot. Since the instruments analyze only a small area on the Sun, a great part of the incoming light needs to be deflected to prevent it from damaging the instruments. At this position a HRW (Heat Rejection Wedge) acts as field stop. Two ammonia-filled heatpipes connect the HRW to dedicated radiators to prevent the HRW from overheating, thus avoiding Schlieren build-up, which could cause wavefront deformations.

The HRW has a central hole of about 2.5 mm diameter (about 10 % of the diameter of the Sun's image). Only the light passing through this hole will hit the telescope‘s secondary mirror to be transmitted to the instruments.

The secondary mirror can be actively controlled to µm precision in order to provide diffraction-limited optical performance even in varying environmental conditions (telescope elevation and changing thermal environment due to passing over cloud decks, sea/land/ice).


Figure 17: HRW (Heat Rejection Wedge) and cooling system (image credit: MPS)


Figure 18: SUNRISE telescope with PFI (image credit: MPS)


SuFI (SUNRISE Filter Imager)

SuFI provides images at violet and near ultraviolet wavelengths. The wavelengths sampled by SuFI are 214 nm at 10 nm bandwidth, 300 nm at 5 nm bandwidth, 312 nm at 1.2 nm bandwidth, 388 nm at 0.8 nm bandwidth, and 396.8 nm(core of Ca II H) at 0.18 nm bandwidth. A 2048 x 2048 UV-enhanced CCD is employed, with a plate scale of 0.02 arcsec/pixel, on average (the plate scale is slightly wavelength dependent). In order to overcome aberrations due to thermoelastic deformations of the telescope and any remaining seeing, a phase-diversity technique is used: a special optical arrangement in front of the detector delivers a nominally focused image on one half of the detector, while the other half receives an image with a defocus of one wave at 214 nm. Hence, the FOV (Field of View) is 15 x 40 arcsec2. Both halves of the CCD together provide sufficient information for post-facto removal of low-order aberrations from the image.

A cadence of better than an image every 2 s can be achieved, depending on the exposure time, to either take rapid time series at a given wavelength or to switch wavelengths. Thus, the four longer wavelengths can be cycled through within 8 s. Since the exposure time of the 214 nm was typically 30 s even at local noon, owing to residual ozone at and above float altitudes, the cadence was correspondingly lower whenever this wavelength was included. 7)


IMaX (Imaging Magnetograph eXperiment)

The IMaX operates in the Fe I 525.02 nm line (a Zeeman triplet with Landé factor g = 3). Images in polarized light covering 50 x 50 arcsec2 are recorded at a spectral resolution of 85 m Å, normally at four wavelengths within the spectral line and one in the nearby continuum. The full Stokes vector in these five wavelengths at a noise level of 10-3 is obtained in 30 s, which is the typical cadence for most of the observations. The number of wavelength points (between 3 and 12) and of polarization states can be varied to obtain a higher cadence, or a better rendering of the line profile shape. 8)

The spectral resolution and sampling is achieved by using a thermally stabilized tunable solid-state Fabry–Pérot etalon in double pass together with a narrowband, prefilter with an FWHM of 0.1 nm.

Polarization states are isolated with the help of two nematic liquid-crystal modulators operated at a frequency of 4 Hz,which are switched between four states for full Stokes vector polarimetry. A dual-beam approach is taken, with two synchronized 1k x 1k CCD cameras. After every observing run, a plate is inserted into the light path in front of one of the cameras in order to obtain phase-diversity information for post-facto reconstruction.

The IMaX optical power system is all refractive; the mirrors of the system are used only for folding and packaging (Figure 19). Three mirrors in total are needed. The optical interface with SUNRISE is the ISLiD focal plane F4 .At F4, a field stop confines the IMaX FOV and prevents unwanted light from entering the instrument and generating parasitic light. This field stop admits two focus positions corresponding to the air and vacuum focus positions, which differ by 1.64 mm in IMaX (the latter being displaced towards the instrument). This shift between the air and vacuum conditions was both predicted by the optical design and checked during the vacuum tests. Next, the pre-filter and LCVRs mounting follow at a distance from F4 of about 30 mm (where the focal depth is ±0.66 mm; near the CCDs it becomes ±2.1 mm). After polarization modulation the beam goes through a collimator system consisting of a doublet and two lenses. The focal length of this collimator is 567.36 mm and its F-number is 25, matching the SUNRISE Telescope and ISLiD F-number. The etalon works in a collimated space producing a blueshift over the FOV. The collimated setup for the etalon requires special care with the maximum angle of incidence on it. The focal length of the collimator guarantees that the angle of incidence on the etalon does not exceed 0.44◦, which produces a wavelength drift of 28 mÅ (over the largest circular FOV).

The diameter of the resulting collimated beam is 25.6 mm, representing the actual area of the etalon used in each pass. In the collimated space between the two folding mirrors, we locate an aperture stop to adequately set the entrance and exit pupils of the system. The beam is finally focused onto the CCDs by a camera consisting of a doublet and two lenses, with a camera focal length of 1021 mm producing a final image IMaX F-number of 44.99. The image focal plane of IMaX has been designed to be telecentric and so has the object focal plane at F4. Both sub-systems, the collimator and the camera optics, are telephoto lenses in order to shorten the total length of the system and get reduction ratios of 0.53 and 0.57, respectively. The magnification of IMaX is 1.799 in order to get a final image scale of 0.055 arcsec pixel-1.

The polarizing beamsplitter is commercial and made of NBK7-Schott glass with anti-reflection coatings on all surfaces. In order to perform phase diversity correction on the image, a parallel plate can be inserted in one of the channels in order to defocus the image on this channel (camera 1). This plate is made of fused silica with 27 mm thickness that produces the specified amount of defocusing corresponding to a phase shift of 1 λ at the edge of the pupil. There are no mechanisms inside the instrument for maintaining IMaX in focus during the flight. For this reason, the most common glass used in IMaX is fused silica, which has near zero thermal expansion (and low stress birefringence). To test and verify the optics of IMaX with a Zygo interferometer (at 632.8 nm), the system has been designed for a wavelength range that includes that of the He-Ne laser. IMaX has thus been achromatized for the range 524.9 nm to 632.8 nm. The materials choice has a great impact on both the achromatization and the athermalization of the instrument. Thus, for the IMaX doublets a pair of glass materials that compensate both the chromatic aberration and the sensitivity to changes in temperature was selected.


Figure 19: IMaX optical (top; self-explained) and optomechanical (bottom) design. In the latter, the F4 plane can be distinguished in a brown-reddish color at the lower left corner. Just after it, the pre-filter and LCVRs housing can be seen in purple. After the stray-light baffle (dark grey), the etalon housing is in blue, preceded by the collimator and camera lenses housing in green. The double light path can be discerned through the two holes in the baffle. To the right-hand side of the etalon enclosure, the vacuum HVPS connectors can be seen in white. After the central baffle, the beamsplitter housing (in pink) appears. Behind it, in yellow, is the PD mechanism. In green are the mountings for the two cameras (bluish) and the PEE. Three isostatic mountings (in brown) for integration into the instrument platform complete the design (image credit: Spanish IMaX consortium)

The influence on the PSF (Point Spread Function) of retroreflection on optical surfaces (ghost images) was analyzed in detail. IMaX includes a high variety of optical components and properties which calls for performing a thorough study on the influence of each surface on the final performances of the instrument. Some of the optical elements, like the etalon and the pre-filter, were considered in a simplified way to reduce the time needed for the stray-light simulation, while still maintaining the strength of the signal. Since the coherence length of light going out from the etalon is quite significant (54 mm for a 50 mÅ spectral resolution), ghost beams should be added in amplitude and not only in intensity. The flux obtained in the detector integrated for the total PSF (including ghost images) has been compared with that obtained from the ghost image alone. The main conclusion of our ghost-image studies was the necessity to tilt the etalon TCPE (Thermally Controlled Pressurized Enclosures) by an angle of 0.36º around the horizontal axis perpendicular to the light path in order to get rid of the ghost images produced by its windows.

Instrument electronics: The IMaX electronics are made up of three main blocks: the main electronics (ME), the optical bench electronics, and the harness. The optical bench contains a proximity electronics (PE), the CCD cameras, and the mechanical, thermal, and optical actuators. Everything is manufactured with commercial-grade, either off-the-shelf (COTS) or specifically designed, components. Both the ME and the PE are enclosed in pressurized vessels (MEE and PEE, respectively) that provide ground pressure conditions during the flight. The harness envelopes are chosen to minimize the outgassing impact in the quasi-vacuum flight conditions. This section is devoted to explaining the various functionalities and the overall design considerations of these three blocks, which are physically and functionally distinct. A block diagram of the IMaX electronics is shown in Figure 20.


Figure 20: IMaX electronics block diagram (image credit: Spanish IMaX consortium)


Figure 21: Top: IMaX optical enclosure just after integration and verification at INTA. The top cover (a radiator) is not in place yet. Labels correspond to the pre-filter and LCVR mounting, the LiNbO3 etalon enclosure, the polarizing beamsplitter, the CCDs, the heat pipes, and the proximity electronics enclosure in the lower left. Bottom: IMaX main electronic enclosure. Labels correspond to the control computer and real-time electronics (DPU), the DC/DC converters, the heat pipes, the fan, the interface board, and the thermal controllers (image credit: Spanish IMaX consortium)

IMaX has been designed, built, and calibrated by a consortium of four institutions in Spain. The Instituto de Astrofísica de Canarias (IAC, Tenerife) is the leading institution and has been in charge of the conceptual design, management of the consortium, and the ground-based control software. The Instituto de Astrofísica de Andalucía (IAA – CSIC, Granada) has been in charge of all the electronic aspects of the instrument. These include the detectors, real-time electronics, control computer and software, power supply and distribution, and harness. The Instituto Nacional de Técnica Aeroespacial (INTA, Torrejón de Ardoz) has developed the optical definition of the instrument, its optomechanical concept, and the thermal sub-system. The Grupo de Astronomía y Ciencias del Espacio (GACE, Valencia) has provided the three enclosures (two of them pressurized) used by the instrument. The AIV (Assembly, Integration and Verification) of the instrument (and various subsystems) has been performed at INTA facilities prior to the submission of the instrument to MPS.


ISLiD (Image Stabilization and Light Distribution)

The ISLiD unit allows simultaneous observations with the two science instruments by distributing the radiation according to wavelength (200–400 nm to SuFI; 525 nm to IMaX; 500 nm to CWS), while preserving diffraction-limited performance as well as polarization information. ISLiD contains a rapid, piezo-driven tip-tilt mirror, which stabilizes the image on the science instrument foci by damping all vibrations and motions acting at frequencies up to 60 Hz. This mirror is controlled by CWS, which is optically stimulated by ISLiD. 9)

ISLiD is a complex optical instrument that has to simultaneously fulfill very different tasks. Firstly, it has to stabilize the incoming beam from the telescope to provide stable output images for the science instruments and the CWS. Secondly, it must divide the incoming light into the wavelength bands required by the instruments, and thirdly, it must feed the instruments with the corresponding wavelengths by providing optical interfaces at given positions and with given directions, plate scale, and pupil position. Since SUNRISE is a high resolution observatory, two requirements must be satisfied: Firstly, the optical system must allow for diffraction-limited performance under all environmental conditions that might arise during flight. Secondly, the image stabilization system must damp the RPE (Residual Pointing Error) to a value which does not degrade the high intrinsic image sharpness by image smear.

It is evident that the different instruments that are fed by ISLiD have individual requirements, which makes it necessary to sort the instruments in a priority list - not in the sense of ranking their scientific importance for the mission, but rather from the point of view of their intrinsic challenges which they impose on the optical design of ISLiD.

Figure 22 shows schematically the SUNRISE optical system and lists the key optical requirements of the individual units from the ISLiD point of view. The most severe requirement of each category is highlighted: the largest simultaneous field of view, the broadest wavelength band, the shortest absolute wavelength, and the fastest working aperture given as effective F-number. Not listed here are any further requirements such as beam direction, image position, or pupil position. Since all instruments were designed to operate at their individual diffraction limit, which is set by the telescope aperture and the specific working wavelength of the instrument, the most critical instrument path is imposed by the SuFI instrument, which aims at a diffraction limit at 214 nm. Also the wavelength bandwidth of SuFI is by far the broadest, spanning almost a factor of two. The SuFI science focus needs a five times larger plate scale than the telescope exit focus (in order to allow for three-pixel oversampling of the point spread function at 214 nm), and therefore a magnification of five is needed from telescope F2 to SuFI science focus. It is evident that diffraction-limited performance at these short wavelengths can only be achieved if the number of optical surfaces is kept to a minimum.

To this end SuFI is completely incorporated into ISLiD, making the SuFI image path an intrinsic part for all instruments. All instrument feeds make use of the SuFI path, which is designed to be diffraction limited at 214 nm, and is therefore uncritical for the longer wavelengths. The separation of the UV light is done after the beam stabilization, and close to SuFI focus, mitigating the disturbing influence of the beam splitter plates, which must be expected to be bent out of shape by their complex coatings. After the separation of the UV light, each of the remaining instruments is fed by a dedicated optical feed providing the required magnification, image, and pupil position. Since CWS and IMaX are operating in the visible range and can be considered as monochromatic, this can best be achieved by dedicated lens optics.

In order to achieve diffraction limited performance of the SuFI path, the following strategy was employed: First, the number of surfaces needed to reimage (magnify) telescope F2 onto SuFI science focus must be minimized. Further, the working FOV of SuFI/ISLiD must be determined by the largest instrument field (50'' x 50'') only and not by the very large telescope exit field of 180 arcsec, which is only required for image motion compensation.

This can be achieved by stabilizing the image before the light enters the reimaging optics: A field lens near the telescope focus images the telescope aperture onto a plane folding mirror, which is actuated by a fast piezo system and acts as fast tip-tilt mirror. Behind this mirror, we only have to consider the reduced FOV of 50 arcsec in the design of the imaging optics. With that, the remaining key requirements for the reimager are: A FOV of 50 arcsec, a working F-number of 24.2 in the incoming path, and an effective F-number of 121 in the exit path. This allows for a design with two spherical mirrors only, which is optimum in terms of design stability: The spherical mirrors are well polishable and not very sensitive to alignment errors, a fact that is of particular importance for a balloon mission, where both, gravity load and thermoelastic deformations change with telescope elevation.


Figure 22: Scheme of the SUNRISE optical system with the relevant optical key requirements form the point of view of ISLiD. Requirements driving the ISLiD design are underlined (image credit: MPS)


Figure 23: Optical design of SuFI: (a) Principle of a Schwarzschild microscope. The magnification of the sketched example design is 5, as in the SuFI instrument. (b) Sketch of the off-axis decentered pupil configuration. The system is stopped down to an entrance F-number of 24.2, yielding an exiting F/121 beam. (c) Folded configuration of the Schwarzschild system. A Lithosil field lens near the object plane images the telescope pupil on the folding mirror, which can now be used as a fast tip-tilt mirror. It compensates image motion in the object plane providing a fixed scene in image plane. (d) Final optical design of SuFI. The exiting beam of the Schwarzschild microscope is folded by a dichroic beam-splitter plate, which reflects all wavelengths below 450 nm. The folded beam passes a double filter set for wavelength selection before the beam is folded again by a UV enhanced aluminum mirror. The phase-diversity image doubler is placed in front of the science focal plane. It provides a focussed and an out-of-focus image of the object side-by-side on the CCD detector.

Phase-Diversity Capability: Phase diversity is a technique for wavefront measuring and subsequent image restoration. Two images of the same object, which are taken simultaneously at two distinct focal positions (usually one image in best focus, the second defocused such as to create an additional wavefront curvature of typically one wave) allow the wavefront in the exit pupil to be sampled at two distinct focal positions. From these two measurements, the object and the aberrations of the system can be retrieved. While for static objects and aberrations the two images can be taken sequentially, for a dynamic object like the Sun the two images must be taken simultaneously.

Phase-diversity capability was incorporated in SUNRISE to correct for residual aberrations due to thermoelastic deformations of the telescope during flight. To this end, the light path includes a phase-diversity image doubler (Figure 24). It resides directly in front of the focal plane and consists of two plane-parallel Suprasil plates, which are coated with dedicated high-reflectivity and anti-reflective coatings on both sides. The second plate is coated on one half of its area with a 50/50 beam-splitter coating, that transmits half of the incoming radiation, while reflecting the other half. This leads to two images lying side-by-side on the focal plane: one image that has passed both glass plates without any reflection, and another image, which has undergone two additional reflections, the first at the beam-splitter plate, the second at the backside of the first plate. The distance and the angle of the plates are chosen such that the path difference between the images is 28.15 mm, while the lateral separation is 12.5 mm, corresponding to half the CCD size. For the F/121 beam this path difference creates a defocus of approximately one wave at 214 nm and half a wave at 388 nm.


Figure 24: Design of the phase-diversity image doubler. A field stop in front of the two glass plates selects the useable field. Half of the light is transmitted through both plates directly and forms the in-focus image, while the other half is optically delayed by two reflections at the specially coated inner surfaces. The path difference is 28.15 mm, corresponding to an additional wavefront curvature of approx. half a wave at 388 nm and approx. one wave at 214 nm. The two images are laterally separated by 12.5 mm, half the size of the CCD (image credit: MPS)


CWS (Correlating Wavefront Sensor)

SUNRISE feeds two scientific instruments: SuFI (SUNRISE Filter Imager) , an UV-imager, and IMaX (Imaging Magnetograph Experiment), a compact filter based magnetograph. In addition, the CWS (Correlating Wave-Front Sensor) is fed by SUNRISE to measure the image shift and possible low-order aberrations (focus and coma), which are due to misalignment of M1 or M2 caused by thermoelastic deformations of the telescope during flight. 10)

CWS consisted of a six-element Shack –Hartmann WFS (Wave-Front Sensor), a fast tip-tilt mirror for the compensation of image motion, and an active telescope secondary mirror for focus correction. The CWS delivered a stabilized image with a precision of 0.04 arcsec (rms), whenever the coarse pointing was better than ±45 arcsec peak-to-peak. The automatic focus adjustment maintained a focus stability of 0.01 waves in the focal plane of the CWS.

For tip-tilt correction, a pupil image of appropriate size is needed where the tip-tilt mirror can be placed. The mirror size determines the necessary mirror stroke (the smaller the mirror the larger the stroke). On the other hand a larger pupil means a longer focal length. The optical solution which was chosen is a field lens which provides a 35 mm pupil image. Here the tip-tilt mirror M105 (notation compatible with official SUNRISE documentation) was placed. This mirror is part of the ISLiD (Image Stabilization and Light Distribution) system (Ref. 9). While maintaining the diffraction-limited image quality, ISLiD provides the following:

• the desired f/121 image for the UV channel (SuFI)

• an f/25 image for IMaX (visible)

• an f/25 image to the CWS (visible)

• a possibility to feed a third instrument (near-IR spectrograph, planned for the second flight).

The main functions of ISLiD are shown schematically in Figure 25. For this purpose, a modified Schwarzschild system is used to magnify the image for SuFI. Four relay lenses demagnify the image and shift the pupil so that IMaX and the CWS are illuminated with an f/25 beam, with a pupil at infinity. An additional correction element accounts for the beam splitter not used in collimated light.


Figure 25: Schematic of the ISLiD system (image credit: MPS)

The CWS is a Shack –Hartmann (SH) type wave-front sensor working at 500 nm. Its optical design is shown in Figure 26. It picks up the image at the entrance focus F6, which also defines the system focus of SUNRISE. Differential foci between the science instruments and the CWS are not corrected by the CWS and have been avoided by a careful system alignment. A field stop with a square area of 1.54 x 1.54 mm2 limits the field of view to 13 arcsec. Coll1 (f = 100 mm) images a 4.6 mm pupil on a hexagonal lenslet array (LLA, f = 81 mm) which consists of seven hexagonal micro-lenses.


Figure 26: Schematic of the wave-front sensor of SUNRISE. The field stop coincides with a focal plane of the telescope (image credit: MPS)

The geometry of the lenslet array is shown in Figure 27. The central one is obscured by the shadow of the telescope's secondary. A relay lens system consisting of a collimator Coll2 (f = 160 mm) and an imaging lens (f = 90 mm) picks up the focal plane of the LLA and images the subfields on the wave-front sensor camera. The image scale is chosen so that 1 arcsec corresponds to five pixels in the camera plane. The left panel of Figure 15 in Section 6 shows the six images on the high-speed camera. It is an original flight picture taken as a screen copy from the ground control system. Solar granulation is clearly visible at high contrast and good image sharpness.


Figure 27: Illumination pattern of the CWS lenslet array. The image of the 1 m entrance pupil provides a homogeneous illumination of the six peripheral micro-lenses, except for the (small) influence of the spiders. The central lenslet is obscured by the secondary mirror and is not used (image credit: MPS)

Wave-Front Correction: An overview of the CWS architecture is given in Figure 28 where black lines are communication paths that are typically labeled with the interface used. The upper region shows the light path from the entrance pupil (EP) of the telescope to the camera. CWS electronics components are shown as yellow boxes. In addition, the figure shows the pointing system (PS, green box) provided by the High Altitude Observatory in Boulder, U.S., and the MTC (Main Telescope Controller, pink box) delivered by the manufacturer of the telescope, Kayser Threde in Munich, Germany. The actual wave-front correction is controlled by a computer named CW-AO depicted in the right part of the overview.


Figure 28: Schematic of the SUNRISE CWS hardware and communication lines. The CWS electronics components are shown as yellow boxes, PS denotes the gondola pointing system and MTC the main telescope controller (image credit: MPS)

Hardware: The CWS electronics consists of two main parts. The MEU (Main Electronics Unit) contains the control processor and the processor for the wave-front correction, as well as all communication and thermal control hardware. The MEU is located outside the telescope, in the shadow of the solar panels, in order to minimize any thermal disturbance to the light path, and in order to avoid external heating by sunlight. The PEB (Proximity Electronics Box) drives the high-speed tip-tilt mirror and is therefore located close to that mirror.

In order to qualify the hardware for the near vacuum conditions of the stratosphere, the project tested the complete CWS (Figure 29) and its electronics at INTA (Instituto Nacional de Técnica Aeroespacial), Madrid, Spain, at pressures down to 100 – 200 Pa and temperatures between +40º and -50ºC. No problems were found during these tests.


Figure 29: The CWS, mounted on its light-weight carbon fiber sandwich plate. On the sides, two fixation points of the isostatic mounting are visible(image credit: MPS)

Thermal and Environmental Design: At flight altitude the atmospheric pressure is about 5 mbar. To dissipate the power effectively by thermal radiation, the cooling elements have an allocated radiation area, and the cooling design has to ensure that the heat is transmitted to that area. For the CW-electronics-unit the project decided to work without a pressurized box, and to dissipate the power only by conduction and radiation. All electronics were designed for low power consumption and small power density to decrease or even avoid any hot spots. The advantages of a conduction cooled electronics unit in contrast to a pressurized box are reduced weight and an uncomplicated temperature control with heaters and radiation coolers. Furthermore, a pressurized box and electronics designed to work under such pressure pose the risk of failure due to loss of air pressure. To protect all components from environmental influences, we applied an electrical isolation and hydrophobic coating to all boards. The cooling elements had a Lord Aeroglaze A 276 coating with high emissivity and low absorption. The power consumption of the CW electronics unit in the operating mode is nearly 70 W, 45 of which are used by the real time computer doing the wave-front correction.

The rapid fluctuations in the temperature profile in Figure 30 show the varying power consumption during the operating and non-operating phases. The low-frequency pattern shows temperature changes between day and night. The decreasing temperature at the beginning results from the ascent of the balloon and the passage through cold atmospheric layers. The noticeable peak on June 11, 2009 coincides with the passage of the gondola over Greenland. The high albedo of the ice led to an increase of the temperature. One of the components of the CWS that required special treatment in order to operate properly at float altitude, was the high-speed camera. We designed a cooling element that connected the hot spot in the interior to the walls of the housing.


Figure 30: Temperature profile of the Main Electronics Box during the balloon flight. The slow variations correspond to temperature changes caused by the varying elevation of the Sun. The strong peak on 11 June occurred during the passage over Greenland, where sunlight was reflected toward the instruments (image credit: MPS)

MEU (Main Electronics Unit): The CWS main electronics unit (Figure 31) contains three different components. The CWControl (a.k.a. CW-Com) with its A9M9750 module based on a NetSilicon's, 200 MHz, NS9750 microcontroller is a compact low energy solution from DIGI. It dissipates nearly 1.7 W in operating mode. The CW-Control includes also the clients of the CoSM-Bus. This bus handles the control data of the motion units. The CW-AO is the real time computer that handles the wavefront correction control loop. It consists of the Rhino DX Board, a conduction-cooled VME single board computer with two 1 GHz Motorola 7457 processors, made by Curtiss Wright, the frame grabber and an interface board, which is specifically designed for the Rhino DX requirements. Attached to CW-AO by two serial interfaces are the digital-to-analog converters of the tip-tilt mirror. Also serially attached to CW-AO are the mechanisms that allow changing between dark stop, field stop and pinhole at the F2 and F6 foci, and the neutral density filter that keeps the WFS camera illuminated appropriately. The power for all components of the CWS is provided from the CW-Supply, which has five voltage levels with different requirements on current or ripple. The power management is controlled by the CW-Control; it permits start up and power down of the whole system in a given order.


Figure 31: Main electronics box of the CWS. 'CW-Control' denotes the control processor (below the cooling plate), the main power supply is labeled 'CW-Supply', 'CW-AO' is the control loop real-time computer (dual Motorola 7457 CPUs), image credit: MPS

PEB (Proximity Electronics Box): The tip-tilt mirror (M105) is driven by the high-voltage amplifier of the PEB (Figure 32). It requires two variable channels and one fixed voltage. Each variable amplifier channel consists of a DA converter with serial interface, a pre-amplifier and a push-pull power stage. The CW-AO dual processor board transmits the data to the high-voltage amplifier via two RS422 serial interfaces.

The motor controller and the high-voltage amplifier are assembled in the CWS PEB that is cooled by conduction. All electronic parts with power losses higher than 250 mW are treated with thermal filler and thereby coupled to the housing. A radiator outside the PEB removes the heat produced by the CWS proximity box. During the flight, the temperatures in the box were between -10ºC and +20ºC.


Figure 32: High-voltage amplifier for the piezo drive of the tip-tilt mirror. From left to right: serial interface with digital/analog converter, pre-amplifier, main amplifier. The green board at the bottom is a CoSM sensor server used for temperature measurements. The black box is the 100 V power supply for the main amplifier (image credit: MPS)

Tip-Tilt Mirror: For the tip-tilt correction a two axis piezo ceramic actuator from Physik Instrumente was used (P.I. S-330K065). This piezo drive fulfills all requirements for the SUNRISE image stabilization (high acceleration, fast response to set points and small increments). The piezo actuators have a total tilt range of 10 mrad (corresponding to ±45 arcsec in the sky). The minimal step size corresponds to 0.15 µrad (1.7 mas on the sky). The piezo actuators have a very low power dissipation. The Zerodur tip-tilt mirror attached to the actuator is 35 mm in diameter with a thickness of 7.5 mm (Figure 33). The surface flatness (75.9 nm peak-to-valley and 12.6 nm rms after attaching the piezos) is important because of its location close to the pupil image. Therefore, it is necessary to mount the mirror in a way that decouples the optical surface from differential thermal expansion of the mechanical parts and to reduce the mounting force as much as possible. The eigenfrequency of the whole system (mirror actuator and mounting) is about 1900 Hz (see Figure 34). At 1700 Hz closed-loop frequency the drive is operated in a minimum of excitation.


Figure 33: Tip-tilt mirror with the piezo stage (in the polished cylinder) and the mechanical interface to the mount assembly (image credit: MPS)


Figure 34: Excitation of the tip-tilt mirror as a function of the sinusoidal excitation frequency (image credit: MPS)

In-Flight Performance:

After the launch and opening of SUNRISE's aperture door, the CWS was the first instrument to be powered up and put into operation. Figure 35 shows one of the first WFS images (left) and the corresponding correlation functions. The excellent quality of both image and correlation functions (compared to a ground-based solar SH-system) is obvious.


Figure 35: Sub-images and correlations functions of the wave-front sensor during closed-loop tracking. Screen shot taken from the flight-control system (image credit: MPS)

In summary, during its successful flight in 2009, the SUNRISE telescope had by far the highest pointing stability at the science focus ever achieved on a balloon-borne telescope. During its five day flight, it provided scientific data of seeing-free, unprecedented polarimetric accuracy and high resolution observations at UV wavelengths never observed before at similar resolution. After the vibrations problems of the gondola have been solved, the goal of 0.005 arcsec rms pointing stability can be reached with the original six sub-aperture setup of the CWS. If some vibrations remain, the pointing could still be improved to about 0.01 arcsec rms by using the increased bandwidth of the two sub-aperture setup. The SUNRISE team is therefore aiming for a second flight, ideally in 2012 close to the solar maximum.


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The information compiled and edited in this article was provided by Herbert J. Kramer from his documentation of: "Observation of the Earth and Its Environment: Survey of Missions and Sensors" (Springer Verlag) as well as many other sources after the publication of the 4th edition in 2002. - Comments and corrections to this article are always welcome for further updates (

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