It is almost impossible to describe LOS (Doppler) velocities of a prominence strictly, due to that a prominence consists of many threads with different velocities [6, 8]. For optically thin lines, such as Si iv 1394 A ̊ , double Gaussian fitting could give two averaged LOS velocities. For optically thick lines, such as Mg ii k/h and H i Lyman series, the profiles are often centrally reversed due to self-absorption, and effects of non-LTE (departure from the local thermodynamic equilibrium) should be considered [16–19]. Despite these difficulties, we could still obtain some information of LOS velocities from the optically thick lines Reference [20]. proposed that for reversed profiles of Mg ii k/h lines, the Doppler shift of the central minimum (or the maximum for a purely emission profile, called k₃ for the k line, see Figure 2J) correlates strongly with the LOS velocity at the τ = 1 height of the line core (τ represents the optical thickness), and the average Doppler shift of the peaks (called k_2b and k_2r for the blue and red sides of the peaks, respectively) correlates with the LOS velocity at the average τ = 1 height of the peaks (deeper than the formation height of the line core). LOS velocity also affects the asymmetry of the peaks [18]. However, it is often difficult to identify the positions of k_2b, k_2r, and k₃ for complex profiles.

The center of gravity (COG) method can be used to derive the Doppler shifts of purely absorption (photospheric lines [21]) or purely emission lines. When there are multi-velocity components in a spectral profile, the COG method gives a weighted average result. However, the COG method is physically wrong for the emission profiles with central reversals, although this method is sometimes still useful in such case [11]. The COG method is expressed as v D = ∑ i ( I ( v i ) − C ) v i ∑ i ( I ( v i ) − C ) , (1) where v _D is Doppler velocity, I is intensity, C is continuum, and wavelength λ is converted into velocity (v) in the unit of “km s⁻¹” using v i = λ i − λ 0 λ 0 c , (2) in which λ ₀ is the rest wavelength and c is the light speed.

To check the reliability of the COG method, we compare it with the methods proposed by Ref. [20]. Figure 2A is a slit-jaw 2796 A ̊ image during the first scan, and Figures 2B,C show the images of the Mg ii k line spectra along slit positions No.56 (at 11:07:39 UT) and No.62 (at 11:08:44 UT), respectively. Four points along the slit position No.56 and four along the No.62 are chosen (denoted with “a”, “b”, “c”, “d” for different heights in panel (A)), and their Mg ii k line profiles are plotted in Figures 2D-K sequentially. Positions of k₃ (marked with the red plus symbols) and k₂ (marked with the green stars) are identified when there is no much confusion. We also calculate the average shift at half maximum for each profile (marked with the blue × symbols). The results are listed in Table 1. From the images of the spectra in Figure 2B-C, it is intuitive that the top part along the slit position No.56 is slightly blue-shifted and the lower part of the prominence is red-shifted; along the slit position No.62, the top part has multi-velocity components and the lower part is also red-shifted. Among the eight line profiles in Figure 2D-K, the positions of k₃ are only identified in four profiles, and three of them are even questionable (56c, 62c, and 62d). So the LOS velocities derived from k₃ are also questionable. The positions of k_2b and k_2r are determined in three profiles relatively precisely. The average shift at half maximum is calculated for each profile, and the results are consistent with what we see in the image of the spectra. However, this method may not include weak components as in case of the profiles from Positions 56a and 62a (Figures 2D,H). The LOS velocities derived from the COG method have the same signs as those derived from the average shifts of peaks or at half maximum except the 62a profile due to the aforementioned reason. The results of the COG method are also consistent with the images of the spectra, which means that this method is not influenced by occurrence of central reversals for the chosen profiles.

We will see that some Mg ii k line profiles are reversed deeply during the prominence eruption. In most cases, the signs of LOS velocities derived from the COG method are generally consistent with the images of the spectra. We will give quantitative results using Gaussian fitting or directly from the images of the spectra, and use the COG method to derive LOS velocities statistically (not physically for profiles with self-absorption) and qualitatively.

The long filament, when the prominence is seen against the solar disk, extends from nearly solar disk center to beyond the northeastern solar limb on 2015 April 28th. Figure 1A–E show its EUV images, including the eruption process in AIA 304 A ̊ . When seen in AIA 171 and 193 A ̊ (Figure 1D–E), the prominence is thinner than that seen in AIA 304 A ̊ (Figure 1A) due to lower opacities. A part of the prominence is blocked by itself before the prominence eruption (Figure 1A), and LOS is mainly along the prominence axis for IRIS observations. In Figure 1B, the prominence is erupting; a dark region along the filament spine and two bright ribbons on the solar disk can be seen. In Figure 1C, the prominence erupts further, and its spine inclines toward the solar equator. We synthesize a time-distance diagram in AIA 304 A ̊ along the slice AB in Figure 1A, and the result is shown in Figure 1F. We fit the prominence height in Figure 1F with a slow-rise phase (linear function) and a fast-rise phase (exponential function) using the approximation [22, 23]. h ( t ) = c 0 e ( t − t 0 ) / τ + c 1 ( t − t 0 ) + c 2 , (3) where h represents height, t is time, t ₀ is arbitrary, and τ, c ₀, c ₁, c ₂ are parameters obtained by fitting. The obtained initial rise speed is c ₁ = 3.78 km s⁻¹, and the erupting speed at the end of our tracking is ∼ 103 km s⁻¹. Onset of the fast-rise phase is defined by t onset = τ ⁡ ln   c 1 τ c 0 + t 0 , (4) which is calculated to be 13:25:30 UT (the cyan vertical line in Figure 1F).

We have seen in Figure 2 around 11:07 UT that the top part of the raster is slightly blue-shifted and the lower part of the prominence is red-shifted. In this section, we will introduce the spectral evolution, especially the variations of LOS velocities of the erupting prominence in detail. Figure 3 shows IRIS observations around 11:17 UT (the second raster scan), including SJI snapshots (Figure 3 left column), images of the spectra along the dark slit in Figure 3A (middle), and spectral profiles (right) of the Mg ii k line (upper part) and the Si iv line (lower part), respectively. Comparing the prominence images in the two wavelengths, SJI 1400 A ̊ has a narrower bright core with faint prominence edges (Figure 3B) than SJI 2796 A ̊ (Figure 3A), and the Si iv 1394 A ̊ line profiles (Figure 3H–J) are weaker with lower signal-to-noise ratios (S/N) than the Mg ii k line (Figure 3E–G). Mg ii k line profiles are sometimes reversed (Figure 3F–G), while Si iv 1394 A ̊ line profiles can be fitted using a Gaussian profile (red curves in Figure 3H–J). Besides, the Mg ii k line has a larger line width than the Si iv 1394 A ̊ , despite that the latter is formed at a higher temperature. These performances are mainly due to that the Mg ii k line has a larger opacity than the Si iv 1394 A ̊ .

In both of the images of the Mg ii k and Si iv 1394 A ̊ spectra in Figure 3C, D, the blue-shifted top part and red-shifted lower part are visible. We choose three regions in Figure 3A, B from right to left with respect to the prominence axis, and their average spectral profiles are plotted in Figure 3E–J. In these panels, the Doppler velocities derived from the COG method (black fonts) and the full widths at half maximum (FWHMs, blue fonts, in the unit of “km s⁻¹”) are denoted at the left. In Figure 3H–J for the Si iv 1394 A ̊ line profiles, the Doppler velocities from Gaussian fitting are denoted at the right (red fonts). Both Mg ii k and Si iv 1394 A ̊ lines show that the right part of the prominence is blue-shifted, with the Doppler velocity of −8.4 ± 0.7 km s⁻¹ (Figure 3H using single Gaussian fitting); the left part is red-shifted, with the Doppler velocity of 11.7 ± 0.8 km s⁻¹ (Figure 3J).

The prominence spectra vary obviously when the prominence approaches eruption. The layout of Figure 4 is similar to Figure 3 but observed at ∼13:15 UT, about 10 min before the onset of the fast-rise phase. In Figures 4C, a red-shifted (around 60 km s⁻¹) faint component appears, which can be seen in SJI at the prominence edges (Figure 4A, B). The Mg ii k line profiles in Figure 4E–G are red-asymmetry with multi-velocity components, and another red Gaussian profile is necessary to fit the Si iv 1394 A ̊ line profiles in Figure 4H–J. With more velocity components, the line widths get wider, and FWHM of the Mg ii k line is nearly 90 km s⁻¹ (Figure 4F, G), and that of the Si iv 1394 A ̊ line is around 50 km s⁻¹ (Figure 4I). However, the Mg ii k line of the faint component has a narrower width around 30 km s⁻¹ (Figure 4C).

During the fast-rise phase in Figure 5, the faint components changes to be blue-shifted, with two main Doppler velocities shown in the image of the Mg ii k line spectra (Figure 5C): -44 and −27 km s⁻¹. In the image of Si iv 1394 A ̊ spectra in Figure 5D, the bright component is also obviously blue-shifted, and the Gaussian fitting results show that the Doppler velocity is at least −17 km s⁻¹ (Figure 5I). However, the bright component in Mg ii k has no obvious shift (Figure 5C), and the line widths are small (Figure 5E, F). We will see in next section that it is due to the absorption by the faint component. The emission of the box 3 in the SJI 2796 A ̊ (Figure 5A) is mainly from the faint component, which is identified from the weak emission of the Mg ii k line with blue shifts (Figure 5G). The weak emission of the Si iv 1394 A ̊ line from the box 3 (Figure 5J) with the similar blue shift suggests that the Si iv line has a contribution to the brightness of the faint component in SJI 1400 A ̊ .

To analyze the evolution of the prominence LOS velocities in detail, we calculate Doppler images from both the Mg ii k and Si iv 1394 A ̊ lines using the COG method. After the calculations, isolated noises in the Doppler images are further removed.

The obtained Doppler images are shown in Figure 6, where the first row shows AIA 304 A ̊ images and the dashed boxes mark the FOV of the Doppler images in the lower two rows. Initially (the first column in Figure 6), the left part of the prominence is red-shifted and the right part is mainly blue-shifted with a boundary near the prominence axis; the maximum red-shifted velocity is ∼ 30 km s⁻¹, and the maximum blue-shifted velocity is ∼ 20 km s⁻¹ from both Mg ii k and Si iv 1394 A ̊ Doppler images (the difference is < 2 km s⁻¹). Then, red shifts dominate the prominence gradually, and the largest Doppler shifts always occur at the left edge. At around 13:15 (Figure 6B5), there is a largely red-shifted region corresponding with the faint component as shown in Figure 4A, B. About 25 min after the onset of the fast-rise phase (the rightmost column in Figure 6), the erupting prominence is mainly blue-shifted, and the average and maximum LOS velocities in Si iv 1394 A ̊ window are 22 and 47 km s⁻¹, respectively. The blue-shifted velocities in Mg ii k window are smaller, and red shifts can still be seen at the left edge. Note that in the Mg ii k Doppler image in Figure 6B6, the positive values along the filament on the solar disk result from the fact that part of the blue wing of the Mg ii k line, which is radiated from the solar disc, is absorbed by the erupting filament (see Figure 5C).

A detailed view on the evolution of the faint component is shown in Figure 7, where the AIA 304 A ̊ images at different times are shown in the first row, the slit-jaw 2796 A ̊ images are shown in the second row, the third row shows images of the Mg ii k line spectra along the slit positions in Figure 7B, and the remaining two rows show Mg ii k profiles at the positions marked in Figure 7C. Panels C-E share the same abscissa but in different units. At 12:55:39 UT (the first column in Figure 7), the regions with or without the faint components can be differed in both AIA 304 A ̊ and SJI 2796 A ̊ . In Figure 7C1, the lower part of prominence spectrum is not blended by the faint component, which is reversed with two clear peaks (Figure 7(e1)); but the upper part is overlaid with the faint component and the peak at the red wing is absorbed with some weak peaks (Figure 7(d1)). Twenty min later (the second column in Figure 7, the same time as Figure 4), the faint component is significantly red-shifted. From Figure 7C2, we can see that the largest Doppler velocity of red-shifted component is beyond 100 km s⁻¹. The Mg ii k line in Figure 7D2 includes information of both bright and faint components; it also peaks at blue wing and the red wing is absorbed. At 13:34:55 UT, about 10 min after the onset of fast-rise phase (the third column in Figure 7), the sign of Doppler shifts of the faint component is changing: the lower part is mainly red-shifted but the upper part is blue-shifted (Figure 7C3). The Mg ii k line in Figure 7D3 is seriously reversed, and the ratio of line peak to central minimum is 4.5. As a comparison, the ratio of the profile in Figure 7E1 is 1.7. At 13:54:33 UT (the fourth column in Figure 7, the same time as Figure 5), the faint component is blue-shifted, and the blue wing of the bright component is absorbed. When focusing on the AIA 304 A ̊ images, we can see the darkening of the faint region during the eruption process. The expansion of the faint prominence spine is visible in both AIA 304 A ̊ and SJI 2796 A ̊ .

A knowledge about the magnetic configuration and viewing angle of the prominence is helpful for understanding its observational features. The left image in Figure 8 is a AIA 193 A ̊ difference image overlaid with LOS magnetogram contours of ±150 Gauss. The post-eruption loops is left skewed, which suggests that the filament channel has a negative magnetic helicity [24]. Hence, we plot the cartoon of the filament as a flux rope [25] with an exaggerated twist number for a clear view (the right image in Figure 8). Due to the fact that the faint component can absorb Mg ii k radiation from the bright component, the faint component should be in front, which could also be seen in AIA 304 A ̊ images (Figure 7A).

Opposite Doppler shifts observed in a same prominence are generally explained as counter-streaming. In previous observations of counter-streaming, bidirectional flows were seen thread by thread [26], or blue- and red-wing images showed different directional flows [27]. Some of counter-streaming observations could be explained by local motions of plasma, such as small-scale oscillations [28] or magnetic reconnection [29]. In our event, the opposite Doppler shifts are found from the beginning of IRIS observations, and are possibly a property of the polar crown prominence that we studied. A major difference from previous observations is that the opposite Doppler shifts are distributed much regularly, i.e., the left half is red-shifted and the right part is mainly blue-shifted. This phenomenon reveals large-scale flows along the prominence spine, and cannot be explained by the local plasma motions.

Unidirectional flows along the flux rope could also cause opposite Doppler shifts when the LOS is perpendicular to the flux rope axis [24, 30–32]. Despite that the LOS is mainly along the prominence spine in the FOV of the IRIS in our case, we assume that the opposite Doppler shifts are caused by the flows within the spine. In this case, in the cartoon of Figure 8, the flows of the part marked with “bright” should move toward the observer, or the flows marked with “faint” should move away from the observer, or both, to cause the left being red-shifted and the right being blue-shifted with respect to the prominence axis. In the following period, the spine is dominantly red-shifted, but the blue shifts are not enhanced. The different evolutions of the red and blue shifts makes this assumption questionable.

During the fast-rise phase in the third column in Figure 7, the faint component shows opposite Doppler shift along the slit, which is possibly caused by untwisting motion of the erupting prominence [10], which could be seen in AIA 304 A ̊ images by tracking the faint region (Figure 7).

When approaching the fast-rise phase, red shifts dominate the prominence, and the faint component has a large LOS velocity. In the Mg ii k window, the LOS velocity of the bright component is not clear due to the fact that its emission is absorbed partly by the faint component. But from the Si iv window, the results of the double Gaussian fitting (Figure 4H–J) suggest that the Doppler shift of the bright component is not obvious. The red shifts could be explained by the mass drainage during the elevation of the prominence spine, which could accelerate the prominence eruption in return [33]. In our case, if the prominence mass moves toward the prominence footpoint behind the solar disk, the direction of LOS velocity is consistent with the red shifts (see the cartoon in Figure 8). The large Doppler velocity of the faint component suggests that the mass near the spine center moves faster than the mass near the footpoint, which is possibly due to more significant elevation of the prominence spine center, and that the large density near the footpoint may slow down the flows.

During the fast-rise phase, the prominence gets to be blue-shifted, which reveals the overall movement of the erupting prominence away from the solar disk and toward the observer.

Reference [10] did also observe a faint and narrow component with a large Doppler velocity up to 460 km s⁻¹ in an erupting active region prominence. The authors proposed that the two components with different Doppler velocities suggest that the erupting material is distributed in a hollow cone shape. The faint component in our observations share most of features as reported by Reference [10]; despite that the Doppler velocity is relatively small. In addition, there are some phenomena only seen in our observations. Firstly, we could see the faint component in AIA 304 A ̊ images and SJIs directly, hence we observed darkening and expansion of the faint component during the prominence eruption (Figure 7A, B). Besides, we observed that the faint component can absorb the Mg ii k radiation from the bright component. On the basis of these phenomena, we suspect that the Mg ii k line from the faint component has a lower opacity than that from the bright component, and the faint component is composed of low-density and cold plasma due to the expansion of the prominence spine.

The intensity ratio of Mg ii k and h lines is widely used to check the opacity [7, 10, 11, 20]: the ratio is generally in the range of 1–2, and approaches 2 under the optically thin assumption and gets small with the increase of opacity. Figure 9 shows the images of Mg ii k integrated intensity (panels A,B) and k/h ratio (panels C,D) from spectral observations. Continuum intensity is subtracted when calculating the images. The left column in Figure 9 is observed between 11:13 UT and 13:19 UT before the occurrence of the faint component, and the k/h ratio is relatively homogeneous. A region is selected in Figure 9C with yellow box, and the mean k/h ratio is 1.25. The right column in Figure 9 is observed between 13:11 UT and 13:17 UT, and a faint region is seen at the prominence left edge (Figure 9B). The k/h ratio image in Figure 9D shows that the faint component has a larger k/h ratio (1.40) than the bright prominence (mainly between 1.2 and 1.3), which suggests that the faint component has a lower opacity. It should be noticed that the k/h ratios vary for different prominences with different viewing angles Reference [11]. reported the k/h value around 1.4 in the main body of an eruptive prominence in quiet region, but in the erupting prominence studied by Reference [10]; the intensity ratio of the primary bright component varies from 1.4 to 1.9.

Prominence Mg ii k radiation is partly from the emission of local plasma, and partly from the scattering of the chromospheric radiation [7, 34]. In our observations, the faint component can absorb or scatter the light from the bright part behind. Using non-LTE radiative transfer techniques Reference [20], simulated emergent Mg ii k and h lines in solar disk; the authors found that Mg ii k core intensity is weak if the line core forms at a high position due to a low density and three-dimensional scattering, although temperature increases along height in the chromosphere. In our observations, the low emission (Figure 7E2-E4) and deep central reversal (Figure 7D3) of the faint component possibly result from the scattering of low-density plasma. In addition, line width is mainly determined by plasma temperature and micro-turbulence. The narrow Mg ii k line profiles suggest that the faint component is mainly prominence core, whose temperature is lower than the PCTR. Therefore, we propose that the faint component consists of low-density and cold plasma, which appear due to the expansion of the prominence core during the prominence eruption. In this process, the flows move fast along the prominence spine and result in significant Doppler shifts. However, non-LTE modeling is necessary to give a strict explanation on the characteristics of the AIA 304 A ̊ and IRIS Mg ii k observations.