SUPPLEMENTARY INFORMATION

Transcription

1 doi: /nature09925 Analysis Details 1. Observations and data analysis The Hinode XRT and SOT images were processed to Level-1 using the SolarSoftware (SSW) routines xrt_prep and fg_prep respectively. The SOT Ca H-line and Hα images were further processed using an unsharp mask filter with a Gaussian profile of 8-pixels FWHM and a weight of 0.5, and 10 pixels and 0.7, respectively. Both image series were then de-jittered using the first frame as an anchor frame and cross-correlating subsequent images to their previous frames in the Fourier domain. The SSW routine tr_get_disp was used to calculate the cross-correlations. Cubic interpolation was used in the image shifting operations. The Hα images were scaled to the plate scale of the Ca H-line images using cubic interpolation and then shifted to align to the Ca H-line images within ±1 Ca H-line pixel. The Ca H-line plate scale 28 is arcseconds pixel -1, summed by 2 on the camera to give a final plate scale of arcseconds pixel -1. The Hα Doppler images were created by subtracting line center +76 må images from the nearest line center -340 må images in time. The Hα line center was estimated to be at -132 må from nominal filter center position giving Doppler wing offsets of ±208 må from line center and peak Doppler velocity sensitivity of ±9.5 km sec -1. Doppler velocity was calibrated to a band of pixels entirely within the spicule forest with an average velocity of the solar rotational red shift at the limb at N56 latitude (approximately 1020 m sec -1 using the filament-prominence solar rotation profile 29 ). The SOT images shown here and in the supplemental movies have been enhanced above the solar limb by a radial filter that reduces disk and spicule intensities by 80 95% relative to their original values. The AIA images are Level-1.5 data that have been enlarged to the SOT Ca H-line plate scale and then shifted and rotated manually to produce alignments that are within about ±5 Ca H-line pixels or approximately 0.54 arcseconds. The enlarged AIA images are enhanced by a Gaussian unsharp mask filter with 60-pixel FWHM and weight of 0.5. All images except the XRT image have been rotated to place solar gravity in the vertical direction. The plume and bubble intensity measurements were made by first tracing the outlines of each feature manually in every frame of the Ca H-line image sequence and then applying the defined region of interest to the scaled and aligned AIA images. The background emission level was determined for the bubble by averaging the intensity in each channel along a 5-pixel wide box from the right edge of the prominence barb to the far right edge of the field of view. The height of the center of the box was chosen to be just above the majority of the limb-brightening emission as defined in the 15:38 UT image series, i.e. approximately 7.5 Mm above the bottom of the images shown in Fig. 2 in the text. For the plume, the background intensity was determined by averaging the levels in a 5-pixel wide box centered on a line through the centroid of the plume (14.5 Mm above the base of images in Fig. 2 in the text) at 15:38 UT, its widest appearance in the time sequence. The levels were again taken from outside the prominence to the far right edge of the field of view (see Figs. 3c g in text). Contrast values as a function of time (Figs. 3 f and g) were determined by dividing the spatially averaged intensity in the region defining the bubble or plume at a given time by the constant value of background intensity determined at 15:38 UT, with the assumption that the background intensity variation is minimal compared to the variation in the prominence structures. This assumption is qualitatively verified by examining the AIA movies associated with this paper. Error bars on the contrast plots denote the standard deviation of the intensity of the pixels in the bubble or plume structures. Filter ratios were taken using intensities normalized by exposure time (as found in the FITS header variable exptime ) to give units of DN/pix/sec. Background emission was determined as described above. Background levels were subtracted from the bubble and plume levels in the 22-June-2010 data before creation of the filter ratio products. For the 02-July-2010 data, the 171/193 ratio used background-subtracted intensities for both channels while the 171/211 and 193/211 ratios were not background-subtracted prior to use due to the 211 intensity being at background intensity and thus causing meaningless ratios. Uncertainties on the AIA filter ratios were determined by assuming that the uncertainty in a given channel intensity measurement is governed by Poisson statistics and propagating the uncertainties to each ratio. The uncertainty in the temperature determination is given by the intersection of the measured ratios with the calculated AIA temperature response ratio curves (see Section SOM 2 below). 1

2 2. AIA Temperature Response and Filter Ratio Analysis The AIA instrument consists of four telescopes, each with at least two spectral bandpasses, or channels. In the EUV, the spectral channels are defined by multi-layer dielectric mirror coatings that select wavelength regions around significant spectral emission lines in the outer solar atmosphere. Each AIA channel is thus tuned to image plasma in the solar atmosphere at the formation temperature of the significant lines. We calculate the temperature response function of the AIA channels by producing a model of the solar emissivity as a function of temperature and wavelength using the CHIANTI atomic database (v ) 30, and then folding the emissivity with the wavelength response of each telescope, assuming photospheric element abundances 31 and using the ionization balance of Bryans et al. 32 Table S1 shows the significant ions and their spectral lines in the main EUV channels as well as the temperature of peak ionization fraction and peak normalized intensity. The temperature response functions of the channels used in this study are plotted in Figure S1. Note that while each channel contains peaks at a characteristic temperature, all have some sensitivity to material outside of their target temperature range. Ratios of the AIA response curves are calculated in an attempt to determine the temperature range for the emitting plasma under the isothermal assumption. The ratios of the AIA response curves are plotted as a function of temperature in Fig. S2. Table S2 lists the emission levels in the 22-June-2010 bubble and plume and Table S3 lists the level in the 02-July-2010 bubble event. The relevant ratios of these quantities are plotted in Fig. S2. Background levels are not subtracted from the levels before forming the ratios since in some cases (e.g. the 22-June bubble) the levels are below the defined background emission to the right of the prominence. We interpret this as a problem in determining a reliable background level in a prominence observation: because the relatively cool prominence absorbs coronal emission, we can only estimate the coronal background by measuring emission levels away from the prominence; it is not clear that the background levels at these remote locations are the same as those behind the prominence. Figure S2a shows the AIA filter ratio plot with measurements for the prominence bubble in the 22-June-2010 data (averaged over the defined prominence bubble and over time for the period 15:32:15 UT to 15:43:58 UT, the period in which the prominence bubble was well defined and generated the large plume shown in Figure 2 of the paper). Fig. S2b shows the same filter ratio curves with the spatially and temporally averaged ratios for the large plume in the 22-June-2010 data. Fig. S2c shows the same filter ratio curves with the spatially averaged ratios taken from the bright plasma structure in the 02-July-2010 prominence bubble at 10:46:59 UT. We did not temporally average the emission in this structure in order to capture the maximum brightness relative to background. Fig. S3d shows the ratios of the AIA 131 channel to the other EUV channels for the 22-June-2010 bubble. The emission levels in the 131 channel for the entire prominence and surrounding corona are quite low and thus the analyses shown in Fig. S3c uses averages of 8 images to increase the signalto-noise ratio. Note that in cases Fig. S2a c there are two solutions given by the intersection of the measured ratios with the curves, one with log T ~ and one with log T ~6.0. In all of these cases the solutions in the higher temperature regime show more overlap than the solutions in the lower temperature regime. Although this is not a rigorous determination, we take this as evidence that the higher temperature solution is more likely than the cooler solution. In addition the radiative lifetime of the 22-June-2010 plume and the bright structure in the 02-July-2010 are compatible with coronal density values. The emission measure at these density levels is not compatible with the cooler solution. Thus although neither discrimination is rigorous due to the lack of a precise density measurement, we can say with some confidence that it is more likely that the plasma in all structures studies is at the higher temperature, i.e. T ~1 MK. The AIA 131 channel analysis was included in an attempt to resolve the ambiguity in the other channel ratios. However as Fig. S2d shows, the 131 ratios actually add a third very hot log T ~ solution to the problem and are thus not helpful. We include the plot here for completeness only. Channel Ions Wavelength Å Fe XXI I I I Fe XXIII I Mg V Fe XX Log T K I max rel. units

3 I I 171 Fe IX Fe X I Ni XIV I Ni XIV Fe IX 193 Fe XII Fe XII Fe XII Fe XI Fe XI Fe IX Fe XI Fe XXIV Fe IX Fe X S XI Fe IX Fe XII 211 Fe XIV Fe XIII Fe XIII Fe XI Fe XIII Fe XIII O IV Ne V Fe XIII Fe XI Cr IX Ni XI O IV N V Ne V 304 He II He II Si XI Table S1. Significant ion emission lines in each AIA spectral passband used in this study. Maximum line intensities (I max ) are normalized by the strongest line in the passband (I 0 ). Only lines with intensities of about 0.01I 0 or greater are listed in the table. 3

4 Channel Wavelength nm Bubble Bubble Background Plume Plume Background SOT Ca II H-line SOT Hα-200mA AIA AIA AIA AIA AIA Table S2. Emission levels in the 22-June-2010 prominence bubble, plume, and background coronal regions. The values for the bubble and plume are spatially averaged over all pixels indentified as belonging to the structures defined in the SOT Ca H-line image and temporally averaged over all images in which the structures were measured. The AIA values do not include pixels that are within the limb-brightened spicular forest region. The bubble and plume background levels are defined on 5-pixel wide linear slices taken outside the prominence region at the same altitude as the structures (see Fig. 1c in main manuscript for the plume slice definition). Channel Wavelength nm Bubble Structure Background AIA AIA AIA Table S3. Emission and background levels in the 02-July-2010 prominence bubble hot structure. The values for the hot structure are spatially averaged over all pixels indentified as belonging to the structure. The structure is defined manually in the AIA 171 image taken at 10:46:59 UT. The background emission is defined as an average over the pixels to the far right of the prominence in a 5-pixel wide linear cut through the centroid of the hot structure. The location of the cut is shown Fig. 4 in the main text. Note that the figure shows a truncated length of the cut on a sub-field of the image used for the measurement. Figure S1. Temperature response functions of the AIA spectral channels used in this study. 4

5 Figure S2. Filter ratio curves and measurements for the three structures discussed in the text. In all panels, the curves show intensity ratios in various AIA channels as a function of temperature for an isothermal plasma in ionization equilibrium. The dashed horizontal lines show the measurements of each ratio for the structure examined with the thick boxes denoting the uncertainty of the ratio measurement (note that in all cases the 171/193 ratio uncertainty is smaller than the line thickness in the plot). The intersection of the measured ratio uncertainty ranges with the curves defines the possible plasma temperatures in the structure (denoted by the shaded vertical boxes). (a) The 22-June-2010 spatially and temporally averaged prominence bubble intensity ratios and the associated temperature ranges. The overlap of implied temperatures is significant only for the higher values of 6.10 < log T < 6.25 (1.3 < T < 1.8 MK). (b) The 22-June-2010 spatially and temporally averaged plume intensity ratio measurements. Again the overlap is more significant only for the higher temperature values of 6.15 < log T < 6.25 (1.4 < T < 1.8 MK). Note that we disregard the intersection with the small peak at log T ~6.5 in the 171/193 ratio since the other ratios are completely incompatible with this solution. (c) The 02-July spatially averaged intensity ratio measurements for the bright bubble structure at 10:46:59 UT. Here the overlap of the lower temperature solutions is very poor while the higher temperature overlap is again much better. The width of the 193/211 ratio temperature boxes here is wider because the measurement line intersects the curves near inflection points which define wider possible temperature ranges then for the other ratios. The higher temperature overlap region gives 6.0 < log T < 6.25 (1.0 < T < 1.8 MK). (d) Inclusion of the AIA 131 channel in the 22-June-2010 bubble analysis does not clarify the temperature of the bubble. Ratios of the 131 emission with the 171, 193, and 211 channels give three intersection solutions including a very hot solution at log T ~ In this case the overlap of solution is best for the log T ~5.4 solution. The uncertainty in all measurements is at or below the thickness of the dashed lines. 5

6 Supporting Discussion 1. The Rayleigh-Taylor Instability in magnetized solar plasmas We have performed numerical simulations of the RT instability in a magnetized solar plasma 33, using the Kippenhahn- Schlüter model 34 as a representation of the prominence magnetic field configuration. Figure S3 shows a 2D slice in the Y-Z plane (plane perpendicular to the line-of-sight) showing the temporal evolution of the magnetic Rayleigh-Taylor instability in the model. The bubble placed inside the prominence is a buoyant, high temperature tube with a non-uniform temperature in the y-direction. The 3D ideal MHD equations are solved using a 2-step Lax-Wendroff scheme based on the scheme presented by Ugai 35. The non-uniform rise of the tube (resulting from the non-uniform temperature) works as a nonlinear perturbation to the system. This results in a high temperature upflow penetrating into the prominence. As the upflow rises, a vortex is created that pushes the prominence material out of the way. The return flows at the bottom of the vortex work to separate the plume upflow from the bubble, closing the prominence material beneath the bubble, as seen in the observational data. Even though a small, random horizontal velocity field was imposed at the start of the simulation, the horizontal flows from the vortex smooth the perturbations above the low temperature part of the bubble so the Rayleigh-Taylor instability in this area is suppressed. These results can give an explanation of the process behind bursting bubbles, where the global rise of the bubble can be stopped, but a locally hotter buoyant component can act as a nonlinear perturbation, driving through the prominence bubble boundary. The vortex motion this creates allows a resealing of the bubble and hence a repeating of the process as is frequently observed in actual solar prominences. It is worth noting that the downflow seen in panels D and H is accelerated to Alfvénic velocity by reconnection of magnetic field lines between the upflow and downflow components, similar to the method for formation of knots presented in Chae et al. 36. The reconnection that occurs in this simulation is numerical. Figure S3. Panels A D show the velocity field (white arrows) over the density (colored contours) as a function of time in the simulation of a Kippenhahn-Schülter prominence model undergoing R-T instability. The dark region at the bottom of the box in Panel A is the low-density bubble; the blue contours show the overlying prominence density. Panels E H show the corresponding temperature profiles with the hot bubble below the relatively cool prominence. Note that the bubble temperature is non-uniform with the higher temperature region on the left side. This is the region that initiates the R-T instability to form the high-temperature plume seen fully developed in Panel D. 6

7 2. Bubble rise height The bubbles from the 22-June and the 02-July events show remarkably different behavior, with the first halting at 11 Mm where as the second reaches approximately 40 Mm where the bubble bursts sending a large fraction of the bubble through the prominence into the coronal cavity above. To understand this in the context of our hypothesis that the bubbles are due to emerging twisted magnetic flux, we examine how an overlying magnetic field can halt the rise of an emergent flux tube. Following the model from Shibata et al. 37 representing the self similar emergence of a magnetic flux sheet into a highly magnetized atmosphere, and adopting a typical plasma-β* for quiescent prominences of , a photospheric pressure scale height (H) of 200 km and a prominence/corona base height of z cor = 12H we can calculate the maximum height a flux sheet should reach as z max = (1 + 1/β) -1/2 exp (z cor /2H 1/2). Figure S4 shows the predicted height that the flux should reach, that is to say the predicted height of a bubble forming below a prominence. The results show that the height of a bubble should be of the order 10 Mm and that the plasma beta range gives a range in height of a factor of 3 times the smallest height, in agreement with the range of heights observed in several SOT observations. Thus the observations of a bubble rising into an overlying prominence and stopping after a certain distance are consistent with the emergence of a magnetic flux rope being halted by the magnetic field in the prominence. In those cases where the bubble rises completely through the overlying prominence as in the 02-July event, we suppose that the magnetic field in the prominence is weaker and hence the plasma-β is larger resulting in a larger bubble height. * The plasma-β parameter is defined as P g /P B where Pg is the gas pressure and P B is the magnetic energy density given by B 2 /2µ 0 with B the magnetic field strength, and µ 0 the magnetic permeability. Figure S4. Plot of prominence bubble height vs. plasma beta parameter in the overlying prominence. 3. Implications of the finding The discovery of emerging bubbles of coronal plasma below solar prominences has major implications for both the theory of quiescent prominence* formation and maintenance as well as the theory of coronal cavity evolution. We consider prominences to be the visible-light manifestation of the larger coronal cavity magnetic flux rope system. In this view, visible-light prominences are the cooler plasma component draining from the lower regions of the coronal flux rope system. Within the coronal cavity flux rope itself we find coronal-temperature plasma with a variety of dynamics including spiral velocity patterns along the interior magnetic field lines. SDO/AIA observations show that coronal temperature plasma is suspended in the concave field lines at the gravitational bottom of the flux rope. Hinode/SOT observations show that the cooler chromospheric plasma drains in vertical streams from the coronal cavity. These observations are consistent with a radiative cooling mechanism that produces the visible-light prominence from plasma initially contained in the coronal flux rope. This coronal condensation mechanism for prominence formation has been criticized in past studies as untenable due to the high mass loss rate in the visible-light prominence drainage; typical observed drainage rates imply that the mass exceeds the total mass in the coronal cavity flux rope by an order of magnitude. But this argument assumes no counterbalancing injection of mass into the coronal cavity. Here we present the first observations of coronal plasma being directly injected into the coronal cavity from below the prominence. The heating is presumably due to coronal heating mechanisms such as current sheet formation and Alfvén wave dissipation taking place within the prominence bubble. Previous ground-based Hα observations (see ref. 13 of main text) have revealed a large prominence bubble with a bright central core. We speculate that the bright central core is seen in 7

8 that observation because of the wide bandpass Hα filter employed and may be evidence of magnetic reconnection and/or a large density enhancement in the bubble due to wave compression. Hinode/SOT observations may not show the bright core in its bubble observations due to the narrow (120 må) filter used for Hα imaging, or because the ground-based event had a rare temperature enhancement that was hotter than anything observed since. Recent Hinode/SOT observations have shown that the drainage of quiescent prominence plasma takes the form of vertically-oriented, narrow filaments separated by turbulent, narrow upflows of dark (low-density) plasma, suggesting constant two-way exchange of mass and, possibly, also magnetic flux between the top and bottom of the prominence. Additional observations 38 show injection into the coronal cavity along the coronal arcade from intermittent small eruptive events. Combined, these observations should motivate future theoretical studies of these coronal mass injection mechanisms as a possible supply of plasma to the coronal cavity that can subsequently drain into the visible-light prominence, as well as a possible supply of magnetic flux to accumulate in the coronal cavity. More generally our discovery points to a new flow system in the solar outer atmosphere which we term magneto-thermal convection to indicate that the buoyancy is due to both thermal and magnetic energy gradients. Convection in an astrophysical context is defined as an energy transport by the movement of fluid in a gravitational field due to thermal gradients and Lorentz forces. It is well understood in the convection zone below the photosphere where the superadiabatic temperature gradient and high plasma-β conditions create cellular convective flow patterns on a range of scales (most notably the photospheric granulation). However our finding is the first indication of a convective instability in the outer atmosphere of the Sun, a region usually assumed to be characterized by simple flows along static, geometrically simple, magnetic field lines in a low-β condition. The presence of a macroscopic sized prominence with its dense filaments and low-density upflows, all embedded at the base of the even larger-scale coronal cavity, tells us that the magnetic field in this entire part of the corona is rich with structures. This magnetic field on the large scale is characterized with by longevity of the order of days to weeks and a continual vertical exchange of mass, suggestive of a magneto-thermal convective phenomenon in its own right. In the new observations we report here, the emergence of a macroscopic-sized, superheated plasma below the prominence a system of chromospheric 10,000 K plasma reveals this magneto-thermal convective phenomenon in a different form, as the interaction between two macroscopic systems in which significant vertical thermal gradients and Lorentz force result in large-scale fluid movements in the gravitational field. In addition to the thermal buoyancy, we presume that the magnetic field in the emerging flux system is significantly higher than in the overlying prominence thus increasing the bubble energy density and buoyancy relative to the overlying prominence. The prominence magnetic fields resist the rise of the buoyant system to create the observed semi-steady-state prominence bubbles. However perturbations to the boundary sufficient to initiate the R-T instability grow into plumes that transfer the hotter fluid through the prominence into the coronal cavity above. In the case of weak overlying prominence fields, the hot emerging flux system rises almost unimpeded as a largescale bubble directly into the coronal cavity. The existence of emerging flux systems below coronal cavities also addresses the question of how large-scale coronal flux ropes might grow quasi-steadily from relatively low-lying compact structures to extended structures visible as fully elliptical cavities high in the corona and hence to eruptive CMEs. Our finding suggests that the intermittent but continual injection of mass and magnetic flux into the coronal cavity from these emerging flux bubbles leads to a build-up of magnetic buoyancy and helicity in the coronal flux rope. This supports theories of CMEs in which loss of magnetic field confinement occurs as the natural end state to a slow quasi-steady evolution of equilibrium states 39,40. * We distinguish here between quiescent and active region/intermediate prominences since the latter type exhibits distinct dynamics (e.g., no bubbles have ever been observed in active region prominences) and may have different formation and maintenance mechanisms. References 28. Shimizu, T., Katsukawa, Y., Matsuzaki, K., Ichimoto, K., Kano, R., Deluca, E., Lundquist, L. L., Weber, M., Tarbell, T. D., Shine, R., Soma, M., Tsuneta, S., Sakao, T., Minesugi, K. Hinode calibration for Precise Image Co-alignment between SOT and XRT, Pub. Astron. Soc. Japan, 59, SP3, S (2007). 29. Allen s Astrophysical Quantities, 4 th Edition, Cox, A. N., ed., Springer-Verlag, New York (2000). 30. Dere, K. P., Landi, E., Young, P. R., Del Zanna, G., Landini, M., Mason, H. E. CHIANTI - an atomic database for emission lines, Astron. & Astrophys., 498, (2009). 31. Grevesse, N., Sauval, A. J. Standard Solar Composition, Space. Sci. Rev., 85, , (1998). 32. Bryans, P., Landi, E., Savins, P. W. A New Approach to Analyzing Solar Coronal Spectra and Updated Collisional Ionization Equilibrium Calculations. II. Updated Ionization Rate Coefficients, Astrophys. J., 691, (2009). 33. Hillier, A., Isobe, H., Shibata, K., Berger, T. E. Numerical Simulations of the Magnetic Rayleigh-Taylor Instability in the Kippenhahn-Schlüter Prominence Model, Astrophys. J., submitted (2011). 8

9 34. Kippenhahn, R., Schlüter, A. Eine theorie der solaren filamente, Z. Astrophys., 43, (1957). 35. Ugai, M. The evolution of fast reconnection in a three-dimensional current sheet system, Phys. Plasmas, 15, 8, (2008). 36. Chae, J., Ahn, K., Lim, E.-K., Choe, G. S, Sakurai, T. Persistent Horizontal Flows and Magnetic Support of Vertical Threads in a Quiescent Prominence, Astrophys. J., 689, L73 L76 (2008). 37. Shibata, K., Tajima, T., Steinolfson, R. S, Matsumoto, R. Two-dimensional magnetohydrodynamic model of emerging magnetic flux in the solar atmosphere, Astrophys. J., 345, (1989). 38. Innes, D., McIntosh, S., Pietarila, A. STEREO quadrature observations of coronal dimming at the onset of mini-cmes, Astron. & Astrophys., 517, L7 L10 (2010). 39. Zhang, M., Low, B. C. The hydromagnetic nature of coronal mass ejections, Ann. Rev. Astron. Astrophys., 43, (2005). 40. Zhang, M., Flyer, N., Low, B. C. Magnetic field confinement in the corona: the role of magnetic helicity accumulation. Astrophys. J., 644, (2006). Movies S1 S10 Movie S1. Hinode/SOT Ca II nm H-line filtergram movie showing the 22-June-2010 quiescent prominence above the NW solar limb. The images have been rotated to show the solar limb horizontal. The pixel scale of the images is arcseconds per pixel. The full field-of-view (FOV) is 130 x 89 arcseconds (95 x 65 Mm) with the rotated SOT FOV contained within. The time span of the movie is from 15:02 UT to 17:44 UT with a 34-minute gap due to an eclipse by the Earth from 16:06 to 16:40 UT. The average cadence (time between frames) in the movie is 14 sec. Note the dark semi-circular bubble and plume formation in the lower right region of the prominence. The plume is visible from approximately 15:34 to 15:44 UT (6,000 sec). Movie S2. Hinode/SOT Hα nm filtergram movie showing the 22-June-2010 quiescent prominence during the same time period as Movie S1. The exact wavelength is -208 må from the line center, i.e nm. The images have been aligned and scaled to the images in Movie S1. Nearest-neighbor interpolation in time is used to temporally align the time series with the SOT Ca H images. The telescope was focused for the Ca II H-line images and shows some defocus artifacts especially near the eclipse. The bubble and plumes are darker in this wavelength compared to the Ca II H-line movie. Movie S3. Hinode/SOT Hα dopplergram movie showing the 22-June-2010 quiescent prominence during the same time period as Movie S1. The images are formed from the Hα ±208 må filtergrams to show red-shifted (positive) velocity as white and blue-shifted (negative) velocity as black in the images. The velocity range is from +10 to -15 km sec -1. Note that the edges of the dark bubble and the plumes visible in the previous movies are red-shifted indicating that the flows are upward and into the plane of the sky. Movie S4. SDO/AIA 304 channel movie showing the 22-June-2010 quiescent prominence during the same time period as Movie S1. The images have been scaled and aligned to the Hinode/SOT Ca II H-line image scale and view field. Nearest-neighbor interpolation in time is used to temporally align the time series with the SOT Ca H images. The dark elongated structures from the peak of the prominence onto the disk are the spine fields of the filament. Note that the bubble and plumes are not visible in this wavelength due perhaps to optical depth effects. Movie S5. SDO/AIA 171 channel movie showing the 22-June-2010 quiescent prominence during the same time period as Movie S1. The images have been scaled and aligned to the Hinode/SOT Ca II H-line image scale and view field. Nearest-neighbor interpolation in time is used to temporally align the time series with the SOT Ca H images. Note that the bubble is somewhat obscured by the normal limb-brightening seen in optically thin EUV emission lines such as the Fe IX 171 Å line that dominates the AIA 171 channel response. Movie S6. SDO/AIA 193 channel movie showing the 22-June-2010 quiescent prominence during the same time period as Movie S1. The images have been scaled and aligned to the Hinode/SOT Ca II H-line image scale and view field. Nearest-neighbor interpolation in time is used to temporally align the time series with the SOT Ca H images. Absorption in the prominence is clearly stronger than in the 171 channel. The bubble and plume are thus seen in higher contrast in this channel. 9

10 Movie S7. SDO/AIA 211 channel movie showing the 22-June-2010 quiescent prominence during the same time period as Movie S1. The images have been scaled and aligned to the Hinode/SOT Ca II H-line image scale and view field. Nearest-neighbor interpolation in time is used to temporally align the time series with the SOT Ca H images. Absorption in the prominence is clearly stronger than in the 171 channel. The bubble and plume are thus seen in higher contrast in this channel. Movie S8. SDO/AIA 131 channel movie showing the 22-June-2010 quiescent prominence during the same time period as Movie S1. The images have been scaled and aligned to the Hinode/SOT Ca II H-line image scale and view field. Nearest-neighbor interpolation in time is used to temporally align the time series with the SOT Ca H images. Prominence emission in the 131 channel is particularly weak so the movie is created from 8-image averages centered on the times of each nearest SOT Ca H image. Movie S9. Composite movie showing a close-up of the large plume in the 22-June-2010 quiescent prominence in three wavelengths: Hinode/SOT Ca H-line, SDO/AIA 171 and SDO/AIA 211 channels. The bubble and plume are easily seen in emission relative to the prominence material. Movie S10. SDO/AIA 304 and 171 channel movies showing the 02-July-2010 quiescent prominence above the SE limb of the Sun. The images have been rotated to show the solar limb horizontal. Notice the large bubble that rises through the prominence until it collapses and then reforms in the same location to collapse again. During the second collapse, the boundary shows clear classic Rayleigh-Taylor symmetric plume/spike formation. Also note the bright emission structure in the 171 channel that rises just below the dark front and then breaks through the front during the cavity collapse. This is the hot core of the emerging structure that causes the large dark bubble visible in the 304 channel. 10