# Undercover EUV solar jets observed by the Interface Region Imaging
Spectrograph

N.-H. Chen1∗,2 and D. E. Innes2 1Korea Astronomy and Space Science Institute,
Daejeon, South Korea
2Max-Planck-Institut für Sonnensystemforschung, D-37077 Göttingen, Germany

###### Abstract

It is well-known that extreme ultraviolet (EUV) emission emitted at the solar
surface is absorbed by overlying cool plasma. Especially in active regions
dark lanes in EUV images suggest that much of the surface activity is
obscured. Simultaneous observations from the Interface Region Imaging
Spectrograph (IRIS), consisting of UV spectra and slit-jaw images (SJI) give
vital information with sub-arcsecond spatial resolution on the dynamics of
jets not seen in EUV images. We studied a series of small jets from recently
formed bipole pairs beside the trailing spot of active region 11991, which
occurred on 2014 March 5 from 15:02:21 UT to 17:04:07 UT. There were
collimated outflows with bright roots in the SJI 1400 Å (transition region)
and 2796 Å (upper chromosphere) that were mostly not seen in AIA 304 Å
(transition region) and AIA 171 Å (lower corona) images. The Si IV spectra
show strong blue-wing but no red-wing enhancements in the line profiles of the
ejecta for all recurrent jets indicating outward flows without twists. We see
two types of Mg II line profiles produced by the jets spires: reversed and
non-reversed. Mg II lines remain optically thick but turn into optically thin
in the highly Dopper shifted wings.The energy flux contained in each recurrent
jet is estimated using a velocity differential emission measure technique
which measures the emitting power of the plasma as a function of line-of-sight
velocity. We found that all the recurrent jets release similar energy (108 erg
cm-2 s-1 ) toward the corona and the downward component is less than 3%.

Sun: corona, UV radiation, particle emission

##  1 Introduction

Jet-like eruptions, occur in coronal holes, active regions, and the quiet Sun
over a broad range of spatial and temporal scales (Shibata et al., 1992;
Shimojo et al., 1996; Cirtain et al., 2007; Moore et al., 2010; Innes et al.,
2016). They have been observed at wavelengths ranging from hard X-ray (HXR) to
white light, revealing their multi-thermal plasma constitution (Chen et al.,
2013). An inversed-Y or Eiffel-tower structure characterizes their standard
morphology and they are typically seen near cancelling and newly emerging
magnetic fields (Shimojo et al., 1996; Canfield et al., 1996; Cheung et al.,
2015). Therefore they are thought to be initiated by magnetic reconnection
(Shibata et al., 1992, 2007).

Jets are often associated with cool plasma surges (Canfield et al., 1996; Chae
et al., 1999; Jiang et al., 2007; Chifor et al., 2008; Sterling et al., 2015;
Mulay et al., 2016). Canfield et al. (1996) observed blue-shifted H$\alpha$
surges adjacent to X-ray jets, associated with moving magnetic features around
a sunspot. Chae et al. (1999) found recurring extreme ultraviolet (EUV) jets
with correlated H$\alpha$ surges near cancelling magnetic flux in a decaying
active region. Recently, Zeng et al. (2016) reported co-spatial H$\alpha$, EUV
and HXR footpoint emission and well-correlated H$\alpha$ and UV spires in
small-scale chromospheric jets observed by the New Solar Telescope (NST). In
these coordinated observations, X-ray/EUV jets and H$\alpha$ surges are seen
from regions with evolving magnetic fields and are characterized by common
features, such as a point-like brightening at the footpoints and a nearly
collimated spire. These studies have shown the close association between
X-ray/EUV jets and H$\alpha$ surges, although a few observational
discrepancies were also noted, such as a time delay between the H$\alpha$
surges and EUV jets (Jiang et al., 2007). Other cool eruptions, such as Ca II
jets, are notable particularly because they share a similar inversed Y-shaped
scheme as the hot X-ray jets in the corona but are much smaller and occur in
the chromosphere (Nishizuka et al., 2008). Furthermore, most of the UV/EUV
jets have been seen at the edge of H$\alpha$ surges (Canfield et al., 1996;
Chae et al., 1999). This could be due to inhomogenous density or temperature
distributions along the outflows or that the cool and hot plasma were ejected
along adjacent field lines (Yokoyama & Shibata, 1995; Anzer & Heinzel, 2005).
The acceleration process may be similar both in cool and hot eruptions, but
the detailed processes have not been clearly determined.

To investigate the nature of the jet-like eruptions, we use data from the two
solar missions: SDO (Pesnell et al., 2012) and IRIS (De Pontieu et al., 2014)
which provide images, spectra and magnetic field strengths from features in
the photosphere, chromosphere, transition region and corona. The Atmospheric
Imaging Assenbly (AIA; Lemen et al., 2012) on SDO consists of seven EUV and
two UV channels with a 12 s and 24 s cadence and $1\arcsec.5$ spatial
resolution of the entire solar disk. The lines are centered at 94, 131, 171,
193, 211, 304, 335, 1600, and 1700 Å. We use images from the Heliospheric and
and Magnetic Imager (HMI; Schou et al., 2012) on SDO to investigate the
underlying photospheric line-of-sight (LOS) magnetic field and intensity. The
IRIS spectrograph observes in two bands: FUV (1332-1406 Å) and NUV (2783-2834
Å) with a spatial and spectral resolution of 0″.33 and 26 mÅ. The slit-jaw
images (SJI) can be obtained simultaneously with spectra in the wavelength
regions 1330, 1400, 2796, and 2830 Å with a spatial resolution of 0.16″
pixel-1. and a cadence of a few seconds.

The aim of this paper is to describe the initiation and evolution of recurrent
jets from a region with moving magnetic features just outside the sunspot
penumbra. In the following, we present a multi-wavelength analysis of the
recurrent jets by using images from IRIS and AIA and spectra from IRIS. In
section 2, the detailed spatial evolution is investigated from the
simultaneous imaging of IRIS and AIA. The corresponding line profiles and the
energy of the recurrent jets are described in section 3. Finally, the summary
and discussion are given in section 4.

##  2 Observations

###  2.1 Overview of observations

Table 1: IRIS observations in OBS 3860260418.

Duration (UT) | 14:49:52 — 17:45:14
---|---
# Raster | 165
Raster Cad. (s) | 64
# Exp./ Raster | 4
Exp. Time | 15.0
Spectral lines (Å) | C II 1336, Fe XII 1349, O I 1356, Si IV 1403, Mg II k 2796
SJIs (Å)/ cadence (s) / exposure (s) | 1400/39/15 ; 2796/32/15 ; 2832/191/15

The events under study were a series of jets that recurred near the periphery
of the trailing sunspot of AR11991 when it was near the disk center on 2014
March 5. The IRIS slit crossed the jets from east to west in a large, sparse,
4-step raster covering a field-of-view (FOV) $3\arcsec\times 119\arcsec$ in 64
s with $1\arcsec$ step size and exposure time 15 s (Table 1). We concentrate
on the IRIS Si IV spectra, IRIS SJI 1400 and 2796 images, and AIA 171, and 304
Å (hereafter, A171 and A304) observations with the objective of investigating
the behaviour of each recurrent jet. The SJI 1400 Å images are dominated by Si
IV line emission from the transition region and continuum emission from the
low chromosphere. The 2796 Å images, dominated by Mg II k lines, are used to
study the corresponding upper chromospheric behaviour of the jets. As for the
two channels of AIA: A171 is emitted primarily from plasma with a formation
temperature at $6\times 10^{5}$ K so should be good for tracking the coronal
evolution of the jets and A304, formed around the same temperature as the Si
IV lines, is dominated by He II emission. Absorption by overlying neutral
plasma may cause a darkening of the EUV emission from the jet and footpoint.
To coalign these images, first, we scaled the AIA and HMI images to match the
FOV and pixel size of the IRIS SJI images. Second, we coaligned the 2832 Å and
HMI continuum images on prominent features, such as the sunspot. We further
matched the 1600 Å with the HMI continuum images, and also the 2832 Å with the
other SJI bandpass images separately. Then the SJI images can be coaligned to
the rest of co-temporal AIA images. We also examined the LOS HMI magnetograms
for clues on the initiation of the jets.

Here we use the strong Si IV lines and also Mg II lines for the spectral study
and present the line profiles in the jet and loop. The shape of the Si IV line
profile in the jet’s spire often consists of two components: a highly blue-
shifted component, namely the moving outflow, and a slight red-shifted
component, representing the background bulk plasma. To begin with the
characteristics of the line profiles, such as the line width are obtained from
double-Gaussian fits. More results will be discussed in section 3.

###  2.2 Recurrent jets

Figure 1: Context images of recurrent jets in AR 11991 on 2014 March 5. Panels
from left to right show coaligned IRIS 1400 and 2796 Å slit-jaw images, AIA
304 and 171 Å images and 171 Å running-difference images. The jets are labeled
as Jet 1, 2 etc. and the associated magnetic bipoles are overplotted on the
left-most frame in which the red/ green contours mark the positive/ negative
polarity with the 10 % and 30 % of maximum counts ($\pm 500$ G) in the HMI LOS
magnetogram. The white arrow drawn on the image taken at time 16:07:43 UT
points to an abnormal brightening (see text).  Figure 2: HMI LOS magnetograms
during the IRIS observation (time starts from left to right, top to bottom as
shown by the black arrows). The red/ green contours mark the positive (white)/
negative (black) polarity with the 10 % (solid lines) and 30 % (dash lines) of
maximum counts ($\pm 500$ G). The right-most bottom frame gives the position
of 1400 Å jets. Circle A surrounds identified biploes related to the jets and
B surrounds the positive field loop footpoint.

Figure 3: Top: (a) Space-time map of SJI 1400 Å emission along the jet. (b)
Temporal evolution of Si IV 1402 Å spectra. In the SJI images, the jet is
traced via a line-cut along the axis of the spire. Dash lines are the slope of
each jet used to calculate the speed of jets and each is tagged with a speed
in km s-1. In the spectral images, the line intensity is averaged for 4th
exposure over the region of the jet (green box in the bottom plots). The time
of each jet is marked by the blue dash lines. Bottom: Regions of interest in
spectra image. It illustrates our selection of specific pixels for estimating
of LOS jet speed. The green box in the spectral image (left) is selected based
on the positions of the spire shown on the slit-jaw images (right). The red
dashed box shows the FOV covered by the slit in one four-step raster. The blue
box covers the region of loop emission. Note these spectra typically have both
high velocity red and blue wings.

The development of the jets is best studied in the 1400 Å images. Figure 1
shows all the recurrent jets (labeled as Jet 1, 2 etc) from 15:02:21 UT to
17:04:07 UT in the 1400, 2796, 304, 171, and 171 Å running-difference images.
The LOS magnetogram contours are overlaid on the leftmost bottom frame to show
the associated magnetic bipoles. Eight jets were observed coming from the same
location during the three hours observation. Jet 1 was the faintest. All of
them show a clear collimated long spire with bright footpoints, namely a
typical inversed-Y shaped geometry, which is in agreement with the standard
solar jet model (Shimojo et al., 1996). Although the contrast of the 2796 Å
images is lower, the identical bright footpoints are visible in these images
as well. The spire itself in the 2796 Å images is actually dimmer than the
surroundings. They have the same length as the bright ones but are
morphologically wider. The dissimilar appearance is due to the line formation
process. While the strong emission of Si IV line is proportional to the column
density of transition region plasma along the LOS, the Mg II k line, emitted
by cooler plasma, is optically thick and its intensity is not a direct
reflection of the column of Mg II along the line-of-sight. The intensity is
given by the source function at optical depth unity. Since the jets are darker
than the surroundings, the source function and possibly the temperature of the
Mg II plasma is cooler in the jet than in the surrounding. Investigation of
the line profiles is required to understand the Mg II emission. The reason
that most the jets are not visible in the EUV images is probably because the
jet EUV emission was absorbed by overlying cool plasma.

These observations demonstrate the importance of the UV observations for
revealing heated jet plasma below the overlying cold material and for
identifying the jet footpoints. Before the IRIS-era, EUV together with
H$\alpha$ images were often used to diagnose the evolution of solar jets and
the accompanying surges (Chen et al., 2015; Zhang & Ji, 2014; Adams et al.,
2014) so it is interesting to see how much additional information on the jet
morphology is revealed by including the UV images. In the jets described here,
typical features such as the bright footpoints and spire, are only partially
visible in the EUV channels, which is also evidenced in the sequences of
running-difference (RD) A171 images (column 5). The RD images are created from
two consecutive images taken 12 s apart and show the intensity
enhancement/decrease of the observed A171 emission in the eruptions. Only two
of the jets (Jet 1 & 2) show simultaneous enhancements of the EUV with the
1400 Å jets. The other jets show no obvious response (mostly grey) in the EUV,
even at their footpoints. The most likely interpretation is surge plasma
absorbing the EUV as indicated by the darkening in the 2796 Å images. It is
noted that there is also an abnormal intensity enhancement (white arrow) in
Jet 2 resulting from the superposition of an additional jet launched from a
different starting point along the path of the other jets (see movie
aia_2.mp4).

Figure 2 shows the temporal evolution of the photospheric LOS magnetic fields
near the base of the recurrent jets. There were two obvious bipole regions
labelled as A and B (white circles) on the south west of the trailing spot of
AR 11991 (see bottom left image in Figure 1). The positive (negative) flux
concentrations are outlined in the red (green) contours with the 10% (solid),
30% (dot) of the maximum magnetic field strength. In the region A, the
positive flux separated and migrated gradually in a west-south direction as
they emerged. Their migration was terminated when they cancelled with the
negative flux. The recurrent jets were launched from region A which was also
the site of the flux cancellation. Region B was more compact and slightly
decomposed at the edge of the negative pole where the footpoint of the loops
were situated.

To estimate the speed of each jet, we make a distance-time map, shown in
Figure 3(a). The jets are traced via a line-cut along the axis of the spire in
the SJI 1400 Å and appear as dark elongations (indicated by blue dashed lines)
in the intensity-reversed map. The speed of each jet ranged from 50 to 200 km
s-1 which is identical to the speed of the jets observed by TRACE in the EUV
(Chae et al., 1999). We also measure their LOS velocity in the averaged 1402 Å
line intensity map shown in figure 3(b). The Si IV line intensity is averaged
over the fourth exposure over the selected pixels (green box in the lower
right hand panel) for the jet’s spire. As each raster consists four 15 s
exposures (plus CCD read-out time $\sim$1.5 s), the plot has a 64 s cadence.
The Si IV 1402.773 Å was used as the reference wavelength for the measurement
of Doppler shifts. The negative Doppler velocity gives the blue shift, namely
the direction away from the Sun while the positive shows the red shift. Each
jet instance was highly blueshifted ($>$100 km s-1) except for the first jet
which was missed by the slit (in the 4th exposure). The LOS velocity from
spectra and the speed derived from the temporal evolution of SJI images are
roughly equal, so the combination gives a jet speed between 70 and 300 km s-1.
The bottom frame of figure 3 shows an example of a SJI image and co-temporal
1400 Å spectral image.

##  3 Time-dependent Spectroscopy

###  3.1 Spectra of jets

Figure 4 gives an overview of the line profiles seen in Si IV and Mg II lines
at selected times. To study the spectra during the jets, we selected the
pixels where the slit crossed the spire and the loops connecting the footpoint
of the jet to the nearby region B (see previous section). The cyan and purple
lines in figure 4 show the selected pixels and the line profile taken
simultaneously with SJI is shown on the right together with the spectral
image. We use 2796.34 Å and 2803.52 Å as the rest wavelengths of the Mg II k
and h lines respectively for measuring the Doppler shifts. The results of
fitting the Si IV spectra are shown with a double-Gaussian at jet’s spire
after background emission from the continuum has been subtracted. In the Si IV
spectra, every jet (blue solid line) had a highly blue-shifted asymmetric
profile or component and almost no enhancement in the red wing, indicating
that this is a collimated outflow without twists. The dominant blue component
of each spire’s profile extends to 200 km s-1 and the width ranged from 23-48
km s-1, suggesting a strong non-thermal upward outflow. The width of their red
wings is bounded between 3-40 km s-1. During the jets, a single-Gaussian
method is used to fit the narrow line profile at the corresponding loop pixels
(red solid line). The loops’ profiles are slightly red-shifted and some have
intense peak emission which is about a factor two stronger than the spires’
emission.

To better understand the behavior of the Mg II lines during the jets, we
consider the k to h intensity ratio in the line core and wings. The Mg II line
formation was discussed extensively in Leenaarts et al. (2013a, b) with
emphasis on the structure of the lines near the core but the profiles from the
jet spire are far from typical due to flows along the jet (Fig. 4), so it is
difficult to select line core and peaks as done by Leenaarts. Instead we have
measured ratios of the k to h at specific Doppler shift in order to determine
whether the emission is optically thin in the jet. As noted, the intensity in
the jet core is significantly less than in the surroundings. This is because
the jet is above the typical chromosphere and in the core, the line is
optically thick so one only sees emission from the optical depth equals one
layer (e.g. Leenaarts et al. (2013a)). In the wings, the intensities in the
jet are higher than in the surroundings and it is hard to tell if the Mg II
lines from the jet is optically thin or thick. If it is optically thin then
the whole column in the jet would contribute emission on top of the background
which is coming from the temperature minimum. If it is optically thick then
the emission is coming from the optical depth equal one layer in the jet and
the emission from the temperature minimum is not seen.

We therefore plot the ratio of the two lines at 0, $\pm 50$ and $\pm 80$ km
s-1 in the jet and at 0 km s-1 in the neighboring loops. Away from the line
core we plot both with and without background subtraction. The background
values are taken as the intensity 0, $\pm 50$ and $\pm 80$ km s-1 before
jetting. We take the average value of intensity at the different velocities
(both k and h lines) in three selected rasters ($\sim$3 mins) and the standard
deviation is used to determine the minimum/maximum background values used in
the line ratio computations. The ratio is very sensitive to the height of the
lines above the background when the difference is small. Therefore the
accuracy of the ratio for weak emission is much less than for stronger
emission.The error bars on the ratios reflect the difference between ratios
computed with a maximum and minimum background level. The error bar is only
applicable and shown when the line intensity is higher than the background
level. In the core (0 km s-1) of the jet’s spire, the k to h ratio is between
1.2 to 1.5 and reduces slightly (1.1$\sim$1.28) at loop pixels (see figure
5a). In the wings of the jets, the ratio with background ranged from 1.19 to
1.8 at $\pm 50$ km s-1and from 0.7 to 1.5 at $\pm 80$ km s-1. After the
background subtraction, it is bounded by same range at -50 km s-1 but drops
down to less than 1 at +50 km s-1. In the high Doppler velocity jet ($\pm 80$
km s-1), several rise up to 2, indicating that the spires are possibly
optically thin. The relation of the k and h lines intensity (figure 5d)
reveals that Mg II lines at the core and wings of the jet are mostly optically
thick even though the central reversal feature sometimes vanishes (e.g. Jet
6). But it could turn into optically thin in the high Doppler shifted wings.

The Mg II spectra at the loops show the standard line profiles with peaks on
either side of the line core. Most cases have a stronger blueward peak (two
nearly equivalent peaks at jet 6 and 8) for both the h and k lines, signifying
downflowing material above the optical depth equals one height (Leenaarts et
al., 2013b). The peak separation is between 25 and 50 km s-1 which is larger
than the separation width of Mg II peaks in a plage region reported by
Carlsson et al. (2015).

Figure 4: The Si IV (left) and Mg II (right) lines profiles. The spectral
images of Si IV (around 1402 Å) and Mg II (2793-2806 Å) at selected time are
overplotted with the line profiles from the pixel that is marked with the a
cyan (jet) or a purple (loop) bar on the left. The count of each line profile
is shown on the right y-axis with units of DN. The co-temporal SJI images of
Si IV and Mg II are shown on the left of the spectra with the locations of the
selected profiles likewise indicated with the cyan and purple bars. Si IV
(blue, red), Mg II k (yellow) and h (green) line profiles as a function of
velocity shifted away from the line center (eg. 1402.77Å, 2796.34 Å and
2803.52Å respectively). The light green dashed-dotted line is the best-fitting
in double- (jet) and single-Gaussian (loop) methods and the fitting parameters
of line width are given on the top of each subfigure.  Figure 5: Mg II k/h
line intensity ratios at (a)line center (vel.= 0 km s-1), (b)vel.= $\pm 50$ km
s-1and (c)vel.= $\pm 80$ km s-1. (d)The scatter plot of k line intensity vs. h
line intensity at vel.= 0, $\pm 50$, $\pm 80$ km s-1 with the orange dash
lines, ratio of 1 and 2 shown .The green and purple circles with x are the
ratio at the jet and neighboring loops respectively. The filled and open
symbols (circle and diamond) in (b), (c) and (d) illustrate the ratio with and
without background subtraction seperately. Crimson and violet red give the
positive (redward) Doppler velocity. Deep and light blue show the negative
(blueward) Doppler velocity. The background are taken as the intensity value
of the vel. at 0, $\pm 50$ and $\pm 80$ km s-1 before the jet occurred. The
error bars in (b) and (c) are the minimum/ maximum ratio values due to the
background subtraction.

###  3.2 Energetics of the recurrent jets

In this section, we address the energy possibly contained in these recurrent
jets. The analysis method adopted here is the Velocity Differential Emission
Measure (VDEM) introduced by Newton et al. (1995) for flare events and
Winebarger et al. (1999, 2002) for explosive events in SOHO/SUMER
observations. The VDEM gives a measure of the energy of the emitting plasma
moving at the LOS velocity in the observed profiles (for more information, see
Newton et al. (1995) & Winebarger et al. (1999)). It is defined by

$\hbox{VDEM}=n^{2}_{e}G(T_{e})\frac{ds}{dv}$ (1)

where $n_{e}$ is the electron density, $G(T_{e})$ is the temperature, $T_{e}$,
dependent emissivity function of the line, s is the distance along the LOS and
v is the velocity. When deriving the VDEM from the observed spectrum, we
assume that the plasma is moving along the LOS direction with uniform density
and pressure during the exposure. Also the observed emission is assumed to
have been emitted from the plasma at the temperature of the peak in the
emissivity function. Since the propagation direction of the jets here is
oblique and not only radially outward, the true energy flux would be larger.
Before applying the method to IRIS observations, we increased the photon
counts (the minimum required photon counts suggested in Winebarger et al.
(2002) is $>$ 1000 counts) by first summing the line profiles of three
consecutive spatial pixels centered at the position with the maximum width
along the slit (mostly around the jet’s spire). We also took the thermal
broadening of Si IV lines (6.88 km s-1 at $T_{e}=8\times 10^{4}$ K), non-
thermal broadening (decided in each profile) and the instrumental broadening
of IRIS (3.9 km s-1, suggested in De Pontieu et al. (2014)) into account. The
corresponding VDEM function is then produced after deconvolution of the
observed spectrum from which the moments of the velocity can be calculated.
Next, we are able to estimate the kinetic energy flux in the plasma
outflow/infow by using equation (1) in Winebarger et al. (2002).

Table 2 summarizes the results of the estimated energy flux of each jet. Three
directional components have been determined, including the energy flux,
kinetic energy flux and skewed energy flux. The directional energy flux is the
total energy carried by the (in-/out-) flows which consists of the total
thermal enthalpy and the total non-thermal energy flux in the certain
direction. The kinetic energy component represents the energy contained in the
flows moving with an average velocity while the skewed component indicates the
high energy contribution, e.g. the high velocity term of the spectral line
profiles (two ends of red and blue wings). Our calculations reveal that the
average energy flux contributed toward the corona (upward) was $2\times$108
ergs cm-2 s-1 and toward the chromosphere (downward) was $2.2\times$106 ergs
cm-2 s-1. Among the average upward energy flux, 39.6% was from the kinetic
energy flux. The downward kinetic energy was 1% of the average downward energy
flux. It is noted that the kinetic is $<$ 50% of the total upward energy in
each jet, suggesting that more than half of energy goes to thermal energy that
heated the surrounding transition region plasma. The energy in the downward
direction is very small.

Table 2: Directional energy flux

| Energy flux
---
($10^{{}_{6}}$ ergs cm-2 s-1)
J1 | J2 | J3 | J4 | J5 | J6 | J7 | J8
$E_{up}$ | 48.02 | 229.93 | 469.93 | 348.04 | 138.67 | 123.92 | 57.01 | 176.53
$E_{down}$ | 1.80 | 1.98 | 0.16 | 2.96 | 1.01 | 6.92 | 2.20 | 0.39
$Ek_{up}$ |  | 11.23
---
(23.3%)
| 67.77
---
(29.4%)
| 218.59
---
(46.5%)
| 141.93
---
(40.7%)
| 48.75
---
(35.1%)
| 42.68
---
(34.3%)
| 16.70
---
(29.2%)
| 83.43
---
(47.2%)
$Ek_{down}$($10^{{}_{4}}$) |  | 4.31
---
(2.4%)
| 4.60
---
(2.3%)
| 0.017
---
(0.11%)
| 0.42
---
(0.14%)
| 0.44
---
(0.44%)
| 5.60
---
(0.81%)
| 2.47
---
(1.1%)
| 0.068
---
(0.2%)

  * •

1: $E_{up}$, $E_{down}$, $Ek_{up}$ and $Ek_{down}$ are the upward/downward
total energy and upward/downward kinetic enegy separately.

  * •

2: The value of each energy flux is the multiples of 106, except for the
$Ek_{down}$ that is for 104.

  * •

3: The precentage of $Ek_{up}$ to $E_{up}$ is gaven underneath and so does the
downward direction.

##  4 Discussion and summary

We studied the spatial and temporal evolution of recurrent jets both with AIA
and IRIS observations in detail. These jets recurred above the active flux
cancellation regions in the periphery of AR 11991. The spectroscopic studies
of these jets suggest collimated outflows without twists. Although the
formation of Si IV and Mg II lines are not the same, the regular features of a
standard jet were commonly visible in the slit-jaw images. We also found
discrepancies in the morphologies of the jets in the UV transition region and
EUV transition region and corona emission (as shown in figure 1). One
explanation is that pre-existing cool material lay between the jets and
observer which caused the absorption of the EUV at short wavelengths (Innes et
al., 2003; Anzer & Heinzel, 2005; De Pontieu et al., 2009). Another
possibility is that the cool jets/ material were accelerated and ejected
because of the slingshot effect caused by the magnetic reconnection if the
cool plasma was situated around the reconnection sites (Yokoyama & Shibata,
1995; Nishizuka et al., 2008). The coexistence of cool and hot plasma
ejections could also be due to the cooling of previous hot ejections/ material
suggested in previous works (Schmieder et al., 1994; Alexander & Fletcher,
1999; Jiang et al., 2007). For our case here, the dim coronal emission
appeared cospatially and cotemporally in the observed wavelengths which makes
the cooling mechanism unlikely. And most of the cool/hot plasma around the
reconnection sites would be exhausted after first ejection (drop by factor of
2 in the following ejection (Archontis et al., 2010)) so it could be difficult
to supply enough cool plasma in the following ejections within the repetition
period (5-15 mins) of jets 2-8 here. It is worth noting that dim coronal
ejections recurred at the same region of AR 11991 throughout the whole day.
The driving mechanism of such long duration dark/dim ejections still needs
further investigations.

The line formation of Mg II lines involves complicated physics which is
greatly influenced by the inhomogeneous solar stratification. The appearance
of the central reversal behaviour in Mg II lines might be related to the ratio
between collision to radiative de-excitation, the density and the temperature
if a simplified radiative transfer model is assumed (Leenaarts et al., 2013a).
We found some profiles in spire position were non-reversed with more symmetric
profiles, e.g. the core in emission, but some were reversed. It might relate
to the enhancement of the density or temperature in each jet-associated
reconnection process. Some flare observations (Liu et al., 2015) suggest that
the non-reversed profiles can be due to high coronal pressures that could be
attributed to the evaporation flows of jets and the reversed ones is often
associated with strong heating. We found the correlation of k to h line
intensity ratio at the line core and wings remains optically thick, however,
it might become optically thin under the highly Dopper shifted condition.

As noted in Si IV lines spectra (figure 4), the broader blueward component is,
the stronger energy flux would be. Energetics of jets derived in Si IV lines
demonstrated a similar magnitude in the upward energy flux with a downward
energy flux of less than 3%. Our case show the average energy flux of these
jets is two-order-magnitude larger than the explosive events reported in SUMER
observations. The energy of those transition region jets or small explosive
events, however, is not the major source to heat the corona or chromosphere.
It is still significantly evidenced that these undercover solar jets,
contained almost similar energy (or mass) as other observed jets, might be
occulted frequently in current EUV observations.

## References

  * Adams et al. (2014) Adams, M., Sterling, A. C., Moore, R. L., & Gary, G. A. 2014, ApJ, 783, 11
  * Alexander & Fletcher (1999) Alexander, D., & Fletcher, L. 1999, Sol. Phys., 190, 167
  * Anzer & Heinzel (2005) Anzer, U., & Heinzel, P. 2005, ApJ, 622, 714
  * Archontis et al. (2010) Archontis, V., Tsinganos, K., & Gontikakis, C. 2010, A&A, 512, L2
  * Canfield et al. (1996) Canfield, R. C., Reardon, K. P., Leka, K. D., et al. 1996, ApJ, 464, 1016
  * Carlsson et al. (2015) Carlsson, M., Leenaarts, J., & De Pontieu, B. 2015, ApJ, 809, L30
  * Chae et al. (1999) Chae, J., Qiu, J., Wang, H., & Goode, P. R. 1999, ApJ, 513, L75
  * Chen et al. (2015) Chen, J., Su, J., Yin, Z., et al. 2015, ApJ, 815, 71
  * Chen et al. (2013) Chen, N., Ip, W.-H., & Innes, D. 2013, ApJ, 769, 96
  * Cheung et al. (2015) Cheung, M. C. M., De Pontieu, B., Tarbell, T. D., et al. 2015, ApJ, 801, 83
  * Chifor et al. (2008) Chifor, C., Isobe, H., Mason, H. E., et al. 2008, A&A, 491, 279
  * Cirtain et al. (2007) Cirtain, J. W., Golub, L., Lundquist, L., et al. 2007, Science, 318, 1580
  * De Pontieu et al. (2009) De Pontieu, B., Hansteen, V. H., McIntosh, S. W., & Patsourakos, S. 2009, ApJ, 702, 1016
  * De Pontieu et al. (2014) De Pontieu, B., Title, A. M., Lemen, J. R., et al. 2014, Sol. Phys., 289, 2733
  * Innes et al. (2016) Innes, D., Bucik, R., Guo, L.-J., & Nitta, N. 2016, ArXiv e-prints, arXiv:1603.03258
  * Innes et al. (2003) Innes, D. E., McKenzie, D. E., & Wang, T. 2003, Sol. Phys., 217, 247
  * Jiang et al. (2007) Jiang, Y. C., Chen, H. D., Li, K. J., Shen, Y. D., & Yang, L. H. 2007, A&A, 469, 331
  * Leenaarts et al. (2013a) Leenaarts, J., Pereira, T. M. D., Carlsson, M., Uitenbroek, H., & De Pontieu, B. 2013a, ApJ, 772, 89
  * Leenaarts et al. (2013b) —. 2013b, ApJ, 772, 90
  * Lemen et al. (2012) Lemen, J. R., Title, A. M., Akin, D. J., et al. 2012, Sol. Phys., 275, 17
  * Liu et al. (2015) Liu, W., Heinzel, P., Kleint, L., & Kašparová, J. 2015, Sol. Phys., 290, 3525
  * Moore et al. (2010) Moore, R. L., Cirtain, J. W., Sterling, A. C., & Falconer, D. A. 2010, ApJ, 720, 757
  * Mulay et al. (2016) Mulay, S. M., Tripathi, D., Del Zanna, G., & Mason, H. 2016, A&A, 589, A79
  * Newton et al. (1995) Newton, E. K., Emslie, A. G., & Mariska, J. T. 1995, ApJ, 447, 915
  * Nishizuka et al. (2008) Nishizuka, N., Shimizu, M., Nakamura, T., et al. 2008, ApJ, 683, L83
  * Pesnell et al. (2012) Pesnell, W. D., Thompson, B. J., & Chamberlin, P. C. 2012, Sol. Phys., 275, 3
  * Schmieder et al. (1994) Schmieder, B., Golub, L., & Antiochos, S. K. 1994, ApJ, 425, 326
  * Schou et al. (2012) Schou, J., Scherrer, P. H., Bush, R. I., et al. 2012, Sol. Phys., 275, 229
  * Shibata et al. (1992) Shibata, K., Ishido, Y., Acton, L. W., et al. 1992, PASJ, 44, L173
  * Shibata et al. (2007) Shibata, K., Nakamura, T., Matsumoto, T., et al. 2007, Science, 318, 1591
  * Shimojo et al. (1996) Shimojo, M., Hashimoto, S., Shibata, K., et al. 1996, PASJ, 48, 123
  * Sterling et al. (2015) Sterling, A. C., Moore, R. L., Falconer, D. A., & Adams, M. 2015, Nature, 523, 437
  * Winebarger et al. (1999) Winebarger, A. R., Emslie, A. G., Mariska, J. T., & Warren, H. P. 1999, ApJ, 526, 471
  * Winebarger et al. (2002) —. 2002, ApJ, 565, 1298
  * Yokoyama & Shibata (1995) Yokoyama, T., & Shibata, K. 1995, Nature, 375, 42
  * Zeng et al. (2016) Zeng, Z., Chen, B., Ji, H., Goode, P. R., & Cao, W. 2016, ApJ, 819, L3
  * Zhang & Ji (2014) Zhang, Q. M., & Ji, H. S. 2014, A&A, 567, A11