HST/COS OBSERVATIONS OF GALACTIC HIGH-VELOCITY CLOUDS: FOUR ACTIVE GALACTIC NUCLEUS SIGHT LINES THROUGH COMPLEX C^*

Our selected AGN sight lines (Mrk 290, Mrk 817, Mrk 876, PG 1259+593) are the first of many AGN targets, scheduled over three years of COS Guaranteed Time Observations (GTOs) that probe HVCs in Complexes C, A, M, WD, and WB. These HVCs were observed previously with UV spectrographs on board HST and FUSE, including the FUSE survey of high-velocity O vi (Sembach et al. 2003) and the HST/STIS survey of high-velocity Si iii (Shull et al. 2009; CSG09). Previously, Mrk 290 was observed with FUSE and HST/GHRS (Wakker et al. 1999; Gibson et al. 2001; CSG03; CSG07). Mrk 817 was observed with HST/GHRS and FUSE (Gibson et al. 2001; CSG03; Fox et al. 2004, 2006). PG 1259+593 was observed with HST/STIS and FUSE (Richter et al. 2001; CSG03; CSG07; Sembach et al. 2004; Fox et al. 2004, 2006). Mrk 876 was observed with HST/GHRS and FUSE (Gibson et al. 2001; CSG03; CSG07; Fox et al. 2004; Sembach et al. 2003).

Table 1 lists the relevant COS observational parameters of the current COS program. Our targets were observed in both the G130M (1134–1480 Å) and G160M (1400–1796 Å) medium-resolution gratings (R ≈ 18,000, Δv ≈ 17 km s⁻¹). The near-UV-imaging target acquisitions were performed with the MIRRORA/PSA mode; the COS primary science aperture (PSA) (∼25 diameter) yields good centering to maximize throughput and resolving power. Mrk 817 was observed during two 2009 epochs (August 4 and December 28). The first-epoch observations were part of the Early Release Observations (EROs) program (11505, PI: Noll), and the second-epoch observations were part of a COS GTOs program (11524, PI: Green). The first observations were obtained during the Servicing Mission Orbital Verification (SMOV) period before the instrument reached final focus. However, the pre-focus data are largely indistinguishable from that taken after correct focus had been achieved except in the case of the very narrowest of absorption features (Δλ ≲ 0.1 Å), none of which are analyzed in this work. Another symptom unique to the early SMOV data was the "divot and clod" feature which occurred near the blue end of each detector segment; charge resulting from photons hitting one part of the detector appeared in a different part. The high voltage on the detectors was reduced on 2009 August 4, and the instrumental feature disappeared. For the Mrk 817 ERO observations, the divot and clod were carefully fitted with Gaussian profiles and normalized to the flux level in adjacent portions of the spectrum. The observing time was approximately equal at both epochs, but the AGN continuum flux decreased by a factor of 1.4 in the interim. Data from the second epoch were scaled multiplicatively to the flux level of the 2009 August observations. A full discussion of the Mrk 817 data appears in Winter et al. (2011). Targets Mrk 290 and PG 1259+593 were both observed as part of COS/GTO programs 11524 (2009 October 29) and 11541 (2010 April 15), respectively. Mrk 876 was observed (2009 April 8–10) in our COS/GTO program and as a Guest Observer target (11686, PI: Arav). The observations were taken within two days of each other.

After retrieval from the Space Telescope Science Institute (STScI), all data were reduced with the COS calibration pipeline, CALCOS⁵ version 2.11f. The analog nature of the COS micro-channel plate makes it susceptible to temperature changes. Electronically injected pulses (stims) at opposite corners of the detector allow for the tracking and correction of any drift (wavelength zero point) and/or stretch (dispersion solution) of the recorded data location as a function of temperature. Flat fielding, alignment, and co-addition of the processed COS exposures were carried out using IDL routines developed by the COS GTO team specifically for COS far-UV data.⁶ The data were corrected for narrow instrumental features arising from shadows from the ion repeller grid wires. We aligned each exposure by cross-correlating strong ISM features and interpolated the aligned exposures onto a common wavelength scale. Wavelength shifts were typically on the order of a resolution element (∼0.07 Å, 17 km s⁻¹) or less. The co-added flux at each wavelength was taken to be the exposure-weighted mean of flux in each exposure.

To transfer the COS data to the V_LSR scale of the 21 cm data, we first aligned the interstellar absorption features in velocity space. Next we chose between 8 and 12 clearly defined HVC absorption features, found the velocity at which their fluxes were minimized, and took the mean. This mean velocity was subtracted from the velocity of the H i 21 cm HVC emission peak, and the resulting difference was used to shift the COS data to align with the 21 cm data. To quantify the quality of the combined data, we identify line-free continuum regions at various wavelengths, smooth the data by the seven-pixel resolution element, and define (S/N)_res = mean(flux)/stddev(flux). The S/N varies across the wavelength range, but representative values in G130M and G160M data are shown in Table 1.

Previous observations from FUSE and HST/GHRS and HST/STIS (CSG03; CSG07; Fox et al. 2004, 2006) were reported on all four sight lines (Mrk 290, Mrk 817, Mrk 876, PG 1259). We included FUSE data (HVC lines of C iii and O vi) and re-measured O vi λ1031.93, using the velocity range defined by the HVC absorbers measured by COS. With the higher S/N, these spectra provide better definition of the extent of high-velocity UV absorption. Tables 3–6 provide the equivalent widths, W_λ, absorption-line data, and derived column densities. Our error bars include statistical fluctuations in the measurement of equivalent widths and systematic errors arising from variations in the continuum placement and HVC velocity range of integration. Our measurements of equivalent widths vary the continuum placement by one standard deviation in continuum flux and adjust the velocity range of HVC absorption by 10 km s⁻¹. These errors are then added in quadrature to produce 1σ confidence limits on our apparent optical depth (AOD) measurements.

The HST/COS data on P ii λ1152.82 and the N i triplet λ1199.54, 1200.22, 1200.70 (toward Mrk 817, Mrk 876, PG 1259) are much better determined than the FUSE upper limits. Other lines in common include Fe ii λ1143.22, 1144.94. The complex line profiles of S ii and Si ii are greatly improved in quality, particularly the S ii lines toward Mrk 876 and Mrk 290. We are also able to provide reliable measurements of N i, P ii, Al ii, and key high ions (C iv, Si iv, N v).

The 21 cm H i data used here were all obtained with the Robert C. Byrd GBT of the NRAO at an angular resolution of 9'. The spectrum for Mrk 876 is that published by Wakker et al. (2011), the spectrum toward PG 1259+593 was obtained from archival GBT data, and new observations were made for Mrk 817 and Mrk 290. In all cases, the data were reduced, calibrated, and corrected for stray radiation following the method outlined in Blagrave et al. (2010) and A. Boothroyd et al. (2011, in preparation). This technique has been shown to produce 21 cm spectra of HVCs, with errors of a few percent in total $N_{\rm H\,{\mathsc{i}}}$ limited by noise and residual instrumental baseline effects. We fitted a second-order polynomial to emission-free regions of each spectrum and smoothed the spectra to produce the following velocity resolution and rms noise: Mrk 876 (0.81 km s⁻¹, 19 mK); PG 1259+593 (0.64 km s⁻¹, 30 mK); Mrk 817 (0.32 km s⁻¹, 27 mK); and Mrk 290 (0.32 km s⁻¹, 25 mK).

Table 2 presents Gaussian fits to the 21 cm HVC components, including their velocities, widths, and H i column densities. We calculated H i column densities for the HVCs in the optically thin assumption (see the footnote to Table 2), which introduces a negligible error for the weak HVC lines. The free parameters are the line centers (V_LSR) of the Gaussians, and their height (brightness temperature T_b) and FWHM in km s⁻¹. We detected H i emission from all HVCs seen in the COS data, except for the most negative velocity component (V_LSR = −275 km s⁻¹ to −175 km s⁻¹) toward Mrk 290. For this range, we estimate an upper limit from the noise in the 21 cm spectrum integrated over the velocity range of the HVC.

All these directions were observed previously in the 21 cm line by Wakker et al. (2001) using the Effelsberg radio telescope (97 beam). The GBT data (90 beam) have lower noise by factors of two to four, slightly better angular resolution, and better spectral baselines. The higher GBT S/N allows us to fit two components to the HVC toward PG 1259+593, whereas Wakker et al. (2001) fitted only one. In general, the two telescopes give similar values of $N_{\rm H\,{\mathsc{i}}}$, to within the errors, except for the −136 km s⁻¹ component of Mrk 290, where the GBT value is 20% below that from Effelsberg, and the −131 km s⁻¹ component of Mrk 876, where the GBT value lies 32% above. In the latter case, the Effelsberg spectrum appears to have been affected by interference. As noted in the captions to Tables 3–6, the total HVC column densities, $\log N_{\rm {{\rm H\,\mathsc{i}}}}$, used in this paper (GBT) are similar to those from Effelsberg (Eff): Mrk 817 (19.50 ± 0.01 from GBT, 19.51 ± 0.01 from Eff); Mrk 290 (20.05 ± 0.01 from GBT, 20.10 ± 0.02 from Eff); Mrk 876 (19.39 ± 0.02 from GBT, 19.30 ± 0.02 from Eff); and PG 1259 + 593 (19.97 ± 0.02 from GBT, 19.95 ± 0.01 from Eff). The primary differences are in the emission-line profiles.

To illustrate the complexities of identifying an HVC in actual UV absorption-line spectra, it is helpful to show true fluxes versus wavelength before deriving equivalent widths from flux-normalized spectra. The UV spectra of AGNs are rich in interstellar lines of metal ions, but they often contain low-redshift IGM absorbers. The interstellar absorbers include low-velocity gas near the LSR, as well as high-velocity and intermediate-velocity gas.

Figures 4–7 present 12 panel plots showing a standard set of HVC absorbers, ranging from low- to high-ionization states. The velocities of Complex C absorption are shown in pink wash. Along two sight lines (Mrk 290 and Mrk 876), additional HVCs at higher negative velocity are shown in blue wash. We show selected interstellar spectral features for the four AGN sight lines: absorption lines of eight low-ionization species (C ii, N i, O i, S ii, Si ii, Fe ii, Al ii, P ii) and four more highly ionized species (Si iii, Si iv, C iv, N v). We also reanalyzed the O vi data from FUSE, using the velocity ranges of HVC ultraviolet absorption, defined by absorption lines of N i, O i, Si ii, S ii, Fe ii, and Al ii. In many panels, we show multiple transitions of the same species, including the N i triplet (1199.550, 1200.223, 1200.710 Å), the S ii triplet (1250.584, 1253.811, 1259.510 Å), the doublets of C iv (1548.195, 1550.770 Å), Si iv (1393.755, 1402.770 Å), and N v (1238.821, 1242.804 Å), two Si ii lines (1190.416, 1193.290 Å), and two lines from the C ii ground state (1334.532 Å) and C ii* fine-structure state (1335.663 Å). One panel shows lines from O i (1302.168 Å) and Si ii (1304.370 Å), and other panels show lines of Si ii (1526.707 Å) and Fe ii (1143.226, 1144.938 Å, and 1608.451 Å).

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The HVC absorption lines are analyzed to provide absorption equivalent width, W_λ, effective line width (b_width), and column density, N, with oscillator strengths from Morton (2003). The column densities were derived using the "AOD" method (Sembach & Savage 1992), which yields the column density, N_a(v), in a given velocity interval. We then derive the total column density by integrating N_a(v) over the velocity range of the HVC. A few of the lines (Si iii, Al ii) are mildly saturated, but we have checked the column densities from the AOD method with the CoGs (curves of growth) for these sight lines, determined from the low ions (CSG03) for accuracy. We believe they are reasonable estimates; further discussion is given in Section 3.4.

The COS line-spread function (LSF) has broad wings with significant power, as documented on the STScI/COS Web site (Kriss 2011). Our COS data-reduction software accounts for the LSF when fitting Voigt profiles to absorption lines, but the LSF is not included in the AOD method for deriving column densities, N_a(v), in fixed velocity bins. Thus, some of the increased velocity range could be instrumental in nature, rather than due to the physical nature of the gas itself. However, the differences in velocity and O vi column densities are not large, as we discuss in the notes on individual sight lines (Section 3.5).

Figures 8–11 show "stack plots" of the HVC absorbers, aligned in velocity space using normalized continua. Our previous experience with the COS wavelength scale (Osterman et al. 2011) suggests that differential shifts up to 10 km s⁻¹ are possible, relative to the 21 cm emission data. In general, the HVC absorption in the UV metal lines agrees well with the H i emission, with the exception of the Mrk 876 sight line. In that case, we integrate the UV absorption in Complex C over the range V_LSR = −170 km s⁻¹ to −100 km s⁻¹, whereas the H i emission appears to have a dip at 〈V_LSR〉 = −160 km s⁻¹, separating Complex C from higher-velocity absorption centered at 〈V_LSR〉 = −190 km s⁻¹ (see Figures 6 and 10).

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From the normalized COS absorption-line data, we derive equivalent widths, line widths, and column densities, which are displayed in Tables 2–5. Details on some of the judgements made in determining column densities are given in Section 3.4. Tables 7–10 compare previous measurements of ion column densities, log N(X_i), together with our adopted values. We also show the ion ratios referenced to $\log N_{\rm {{\rm H\,\mathsc{i}}}}$ and the corresponding elemental abundances.

Table 7. Mrk 817: Summary of Column Densities^a and Abundances

Species	CSG-03	Shull-11	Adopted	Abundanceb	Abundanceb
(Xi)	($\log N_{X_i}$)	($\log N_{X_i}$)	($\log N_{X_i}$)	$\log (N_{X_i}/N_{\rm H\,\mathsc{i}})$	[X/H]
O i	15.72+0.24 − 0.16	⩾14.98	15.72+0.24 − 0.16	−3.78 ± 0.24	−0.47 ± 0.24
N i	<14.05	14.06+0.12 − 0.07	14.06+0.12 − 0.07	−5.44 ± 0.10	−1.27 ± 0.10
S ii	14.34+0.05 − 0.05	14.29+0.08 − 0.08	14.29+0.08 − 0.08	−5.21 ± 0.08	−0.60 ± 0.10
Si ii	14.48+0.07 − 0.08	14.47+0.08 − 0.05	14.47+0.08 − 0.05	−5.03 ± 0.08	−0.75 ± 0.10
Fe ii	14.31+0.11 − 0.08	14.23+0.10 − 0.06	14.23+0.10 − 0.06	−5.27 ± 0.10	−0.88 ± 0.15
C ii	...	⩾14.9	⩾14.9	> − 4.6	⩾ − 1.03
Al ii	...	13.23+0.07 − 0.06	13.23+0.07 − 0.06	−6.27 ± 0.07	−0.72 ± 0.07
P ii	...	<12.95	<12.95	< − 6.55	<0.04
C iv	...	13.77+0.14 − 0.09	13.77+0.14 − 0.09	−5.73+0.14 − 0.09
N v	...	13.16+0.10 − 0.13	13.16+0.10 − 0.13	−6.34+0.10 − 0.13
O vi	...	14.05+0.16 − 0.09	14.05+0.16 − 0.09	−5.45+0.16 − 0.09
Si iv	...	13.08+0.13 − 0.12	13.08+0.13 − 0.12	−6.42+0.13 − 0.12
Si iii	...	13.74+0.15 − 0.13	13.74+0.15 − 0.13	−5.76+0.15 − 0.13

Notes. ^aComparison of measured column densities, log N_a(cm⁻²), for HVC at 〈V_LSR〉 = −109 km s⁻¹ (integrated from −190 km s⁻¹ to −70 km s⁻¹). ^bAbundances for neutrals and first ions X_i are given relative to measured $\log N_{\rm H\,\mathsc{i}} = 19.50 \pm 0.01$ (Table 2). Fox et al. (2004) found log N(O vi) = 13.97^+0.10 _{− 0.13} for HVC absorption between −160 and −80 km s⁻¹. Last column gives abundances of elements [X/H] relative to solar abundances (Asplund et al. 2009), with ionization corrections of 0.27 dex (S ii), 0.21 dex (Si ii), and 0.11 dex (Fe ii) subtracted from the first-ion abundances (Section 3.2).

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Table 8. Mrk 290: Summary of Column Densities^a and Abundances

Species	CSG-03	CSG-07	Shull-11	Adopted	Abundanceb	Abundanceb
(Xi)	($\log N_{X_i}$)	($\log N_{X_i}$)	($\log N_{X_i}$)	($\log N_{X_i}$)	$\log (N_{X_i}/N_{\rm H\,\mathsc{i}})$	[X/H]
O i	<16.79	15.75+0.29 − 0.15	14.64+0.03 − 0.02	15.75+0.29 − 0.15	−4.30+0.29 − 0.15	−0.99+0.29 − 0.15
N i	<15.02	14.23+0.17 − 0.20	14.06+0.12 − 0.07	14.06+0.12 − 0.07	−5.99+0.12 − 0.07	−1.82+0.12 − 0.07
S ii	14.29+0.15 − 0.14	14.24+0.16 − 0.20	14.43+0.08 − 0.08	14.43+0.08 − 0.08	−5.62+0.08 − 0.08	−0.78+0.15 − 0.15
Si ii	<14.97	14.93+0.18 − 0.13	>14.40	14.93+0.18 − 0.13	−5.12+0.18 − 0.13	−0.67+0.15 − 0.15
Fe ii	<15.46	14.41+0.23 − 0.21	14.26+0.09 − 0.05	14.26+0.09 − 0.05	−5.79+0.09 − 0.05	−1.31+0.09 − 0.05
Al ii	...	...	13.11+0.12 − 0.09	13.11+0.12 − 0.09	−6.94+0.12 − 0.09	−1.39+0.12 − 0.09
P ii	...	<12.82	12.83+0.22 − 0.18	12.83+0.22 − 0.18	−7.22+0.22 − 0.18	−0.63+0.22 − 0.18
C iv	...	...	13.78+0.14 − 0.09	13.78+0.14 − 0.09	−6.27+0.14 − 0.09
N v	...	<13.49	12.98+0.22 − 0.15	12.98+0.22 − 0.15	−7.07+0.22 − 0.15
O vi	...	14.23+0.04 − 0.04	14.10+0.15 − 0.14	14.10+0.15 − 0.14	−5.95+0.15 − 0.14
Si iv	...	...	12.95+0.12 − 0.12	12.95+0.12 − 0.12	−7.10+0.12 − 0.12
Si iii	...	...	13.15+0.10 − 0.08	13.15+0.10 − 0.08	−6.90+0.10 − 0.08

Notes. ^aComparison of measured column densities, log N_a(cm⁻²), for HVC at 〈V_LSR〉 = −120 km s⁻¹ (integrated from −175 km s⁻¹ to −75 km s⁻¹). ^bAbundances for neutrals and first ions X_i are given relative to measured $\log N_{\rm H\,\mathsc{i}} = 20.05 \pm 0.02$ for all three HVC components (Table 2). Wakker et al. (1999) found log N(S ii) = 14.34^+0.08 _{− 0.11}. The last column gives elemental abundances [X/H] relative to solar values (Asplund et al. 2009) with ionization corrections of 0.04 dex (S ii), 0.04 dex (Si ii), and 0.02 dex (Fe ii) subtracted from the first-ion abundances (Section 3.2).

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Table 9. Mrk 876: Summary Column Densities^a and Abundances

Species	CSG-03	CSG-07	Shull-11	Adopted	Abundanceb	Abundanceb
(Xi)	($\log N_{X_i}$)	($\log N_{X_i}$)	($\log N_{X_i}$)	($\log N_{X_i}$)	$\log (N_{X_i}/N_{\rm H\,\mathsc{i}})$	[X/H]
O i	15.55+0.42 − 0.28	15.26+0.17 − 0.14	>14.74	15.26+0.17 − 0.14	−4.13+0.17 − 0.14	−0.82+0.17 − 0.14
N i	14.20+0.15 − 0.14	13.90+0.03 − 0.04	13.97+0.13 − 0.08	13.97+0.13 − 0.08	−5.42+0.13 − 0.08	−1.25+0.13 − 0.08
S ii	...	<14.34	14.25+0.13 − 0.14	14.25+0.13 − 0.14	−5.14+0.13 − 0.14	−0.58+0.20 − 0.20
Si ii	14.53+0.11 − 0.13	14.56+0.11 − 0.09	14.41+0.16 − 0.09	14.41+0.16 − 0.09	−4.98+0.13 − 0.14	−0.76+0.20 − 0.20
Fe ii	14.41+0.09 − 0.08	14.36+0.07 − 0.07	14.26+0.17 − 0.10	14.26+0.17 − 0.10	−5.13+0.17 − 0.10	−0.77+0.20 − 0.20
C ii	...	...	>14.76	>14.79	> − 4.60	> − 1.03
Al ii	...	13.43+0.31 − 0.19	13.12+0.07 − 0.11	13.12+0.17 − 0.11	−6.27+0.17 − 0.11	−0.72+0.20 − 0.20
P ii	<13.21	...	<12.97	<12.97	< − 6.42	<0.13
C iv	...	13.80+0.04 − 0.03	13.58+0.29 − 0.18	13.58+0.29 − 0.18	−5.81+0.29 − 0.18
N v	...	<13.32	12.85+0.23 − 0.14	12.85+0.23 − 0.14	−6.54+0.23 − 0.14
O vi	...	14.20+0.02 − 0.02	13.99+0.36 − 0.14	13.99+0.36 − 0.14	−5.40+0.36 − 0.14
Si iv	...	13.28+0.03 − 0.02	13.10+0.25 − 0.15	13.10+0.25 − 0.15	−6.29+0.25 − 0.15
Si iii	...	...	13.72+0.24 − 0.14	13.72+0.24 − 0.14	−5.67+0.24 − 0.14

Notes. ^aComparison of measured column densities, log N_a(cm⁻²), for HVC at 〈V_LSR〉 = −133 km s⁻¹ (integrated from −170 km s⁻¹ to −100 km s⁻¹). ^bAbundances for neutrals and first ions X_i are given relative to measured $\log N_{\rm H\,\mathsc{i}} = 19.39 \pm 0.01$ for the −131 km s⁻¹ component (Table 2). Fox et al. (2004) found log N(O vi) = 14.12^+0.11 _{− 0.13} and log N(N v) <13.43. Our previous Si iii survey (Shull et al. 2009) measured log N(Si iii) ⩾13.92 as adopted here. The last column gives elemental abundances [X/H] relative to solar values (Asplund et al. 2009), with ionization corrections of 0.32 dex (S ii), 0.27 dex (Si ii), and 0.14 dex (Fe ii) subtracted from the first-ion abundances (Section 3.2).

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Table 10. PG1259+593: Summary of Column Densities and Abundances^a

Species	CSG-03	Richter-01	Shull-11	Adopted	Abundanceb	Abundanceb
(Xi)	($\log N_{X_i}$)	($\log N_{X_i}$)	($\log N_{X_i}$)	($\log N_{X_i}$)	$\log (N_{X_i}/N_{\rm H\,\mathsc{i}})$	[X/H]
O i	15.75+0.18 − 0.24	15.77+0.37 − 0.31	...	15.85+0.15 − 0.15	−4.12+0.15 − 0.15	−0.81+0.15 − 0.15
N i	14.02+0.19 − 0.12	13.95+0.17 − 0.21	13.85+0.12 − 0.12	13.85+0.12 − 0.12	−6.12+0.12 − 0.12	−1.95+0.12 − 0.12
S ii	14.38+0.12 − 0.11	14.34+0.12 − 0.15	14.35+0.15 − 0.10	14.35+0.15 − 0.10	−5.62+0.15 − 0.10	−0.78+0.15 − 0.10
Si ii	14.67+0.20 − 0.15	14.56+0.28 − 0.27	14.17+0.10 − 0.14	14.17+0.10 − 0.14	−5.80+0.10 − 0.14	−1.35+0.10 − 0.14
Fe ii	14.40+0.17 − 0.08	14.16+0.20 − 0.14	14.11+0.12 − 0.11	14.11+0.12 − 0.11	−5.86+0.12 − 0.11	−1.38+0.12 − 0.11
C ii	...	...	>14.45	>14.45	> − 5.52	> − 1.95
Al ii	13.45+0.17 − 0.16	13.42+0.30 − 0.50	12.89+0.13 − 0.16	12.89+0.13 − 0.16	−7.08+0.13 − 0.16	−1.53+0.13 − 0.16
P ii	<12.96	<13.22	<12.43	<12.43	< − 7.54	< − 0.95
C iv	...	...	13.08+0.33 − 0.15	13.08+0.33 − 0.15	−6.89+0.33 − 0.15
N v	...	...	<12.99	<12.99	< − 6.98
O vi	...	...	13.58+0.18 − 0.12	13.58+0.18 − 0.12	−6.39+0.18 − 0.12
Si iv	...	...	12.64+0.17 − 0.13	12.64+0.17 − 0.13	−7.33+0.17 − 0.13
Si iii	...	...	13.30+0.24 − 0.16	13.30+0.24 − 0.16	−6.67+0.24 − 0.16

Notes. ^aComparison of measured column densities, log N_a(cm⁻²), for HVC at 〈V_LSR〉 = −128 km s⁻¹ (integrated from −170 to −95 km s⁻¹). ^bAbundances for neutrals and first ions X_i are given relative to measured $\log N_{\rm H\,\mathsc{i}} = 19.97 \pm 0.02$ for the HVC components at −126 km s⁻¹ and −129 km s⁻¹ (Table 2). We adopted $\log N_{{\rm O\,\mathsc{i}}} = 15.85 \pm 0.15$ from Sembach et al. (2004). Fox et al. (2004) found log N(C iv) =13.26^+0.04 _{− 0.06}, log N(N v) <12.85, log N(Si iv) =12.73^+0.05 _{− 0.03}, and log N(O vi) =13.71^+0.09 _{− 0.09}. Our previous Si iii survey (Shull et al. 2009) measured log N(Si iii) ⩾13.60 as adopted here. The last column gives elemental abundances [X/H] relative to solar values (Asplund et al. 2009), with ionization corrections of 0.04 dex (S ii), 0.04 dex (Si ii), and 0.02 dex (Fe ii) subtracted from the first-ion abundances (Section 3.2).

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Tables 7–10 list the adopted and abundances, relative to solar values taken from a recent review (Asplund et al. 2009). On a logarithmic scale where hydrogen is 12.00, these solar values are C (8.43), N (7.83), O (8.69), S (7.12), Si (7.51), Fe (7.50), Al (6.45), and P (5.41). The previous Complex C abundance estimates were derived from [O i/H i] and [S ii/H i]. The O i measurements should provide an accurate determination of [O/H], owing to the strong resonant charge-exchange coupling of O i and O ii with H ii and H i. However, in many sight lines, the O i line (λ1302.16) is saturated, and we adopted previous values from weaker lines in the FUSE band. Abundances from other ions such as [S ii/H i], [Si ii/H i], and [Fe ii/H i] require ionization corrections to arrive at the true metallicities of these elements. There is no guarantee that all elements have the same metallicity, owing to the chemical history of the HVC or possible dust depletion of refractory elements (Si, Al, Fe). As shown in our previous work (see Figure 16 of CSG03), these corrections are typically expressed as the logarithmic difference between the ion abundance and elemental metallicity. For example, the difference [S ii/H i]−[S/H] ranges from 0.1 to 0.6 for H i column densities $\log N_{\rm {{\rm H\,\mathsc{i}}}} =$ 19.4–20.1 and HVC physical densities n_H = 0.01–0.1 cm⁻³. The HVC metallicities are usually found from [O i/H i], which requires no ionization correction (CSG03; CSG07), and from [S ii/H i], reduced by a factor that depends on the H i column density.

Our four AGN sight lines fall into two distinct groups: one with low column density (Mrk 876 and Mrk 817 have $\log N_{\rm {{\rm H\,\mathsc{i}}}} \approx 19.4\hbox{--}19.5$) and one with high column density (Mrk 290 and PG 1259+593 have $\log N_{\rm {{\rm H\,\mathsc{i}}}} \approx 20.0\hbox{--}20.1$). The ionization corrections to the metal ions (CSG03) depend on the "photoionization parameter," the ratio of ionizing radiation field to gas density n_H. As seen in Figure 1, the HVC clumps have characteristic angular sizes of 1°–2°, perhaps poorly characterized owing to the 06 beam and 05 sampling grid of the LAB survey. At the 10 kpc distance of Complex C, 1° corresponds to 175 pc. If the H i absorbers have a comparable depth and angular extent, the observed H i column densities correspond to physical densities n_H ≈ 0.05 cm⁻³ (Mrk 876 and Mrk 817) and n_H ≈ 0.2 cm⁻³ (Mrk 290 and PG 1259). From our previous photoionization modeling and observed H i column densities, we adopt the following corrections. For Mrk 290 and PG 1259, the corrections are minor: [S ii/H i] and [Si ii/H i] are reduced by 0.04 (dex) and [Fe ii/H i] is reduced by 0.02 (dex). For Mrk 876, we reduce [S ii/H i] by 0.32, [Si ii/H i] by 0.27, and [Fe ii/H i] by 0.14. For Mrk 817, we reduce [S ii/H i] by 0.27, [Si ii/H i] by 0.21, and [Fe ii/H i] by 0.11. We have not made ionization corrections for O i, N i, Al ii, or P ii. The ionization corrections noted above for S ii, Si ii, and Fe ii have been applied and included in the final column of Tables 7–10.

The Si iii column density can be used as a proxy for ionized gas to infer the column density, N(H ii), of ionized hydrogen that is kinematically associated with the observed H i. As discussed by Shull et al. (2009), the Si iii ion likely includes contributions from both photoionized and collisionally ionized gas (only 16.34 eV needed to produce Si iii). We adopt an ionization fraction $f_{\rm Si\,\mathsc{iii}} = 0.7 \pm 0.2$ characteristic of multiphase conditions. Following the methodology of our Si iii survey with HST/STIS (Shull et al. 2009), we assume a metallicity Z_Si ≈ 0.1 relative to the solar abundance, (Si/H)_☉ = 3.24 × 10⁻⁵, to find

$$\begin{equation} N_{\rm H\,\mathsc{ii}} = (4.4 \times 10^5) N_{\rm Si\,\mathsc{iii}} \left[ \frac{Z_{\rm Si}}{0.1} \right] ^{-1} \left[ \frac{f_{\rm Si\,\mathsc{iii}}}{0.7} \right] ^{-1} \;. \end{equation} \tag{ 1 }$$

For the current four sight lines, we infer ionized column densities, log N(H ii) =18.79 (Mrk 290), 19.24 (PG 1259+593), 19.38 (Mrk 817), and 19.56 (Mrk 876). Compared with the observed H i, these correspond to ionization fractions, N(H ii)/N(H i), ranging from low values of 6%–10% in the high-column HVCs toward Mrk 290 (log N$_{\rm {{\rm H\,\mathsc{i}}}} = 20.05$) and PG 1259+593 (log N$_{\rm {{\rm H\,\mathsc{i}}}} = 19.97$) to much larger values of 80% and 150% toward Mrk 817 (log N$_{\rm {{\rm H\,\mathsc{i}}}} = 19.50$) and Mrk 876 (log N$_{\rm {{\rm H\,\mathsc{i}}}} = 19.39$). Interestingly, the highest ionization fractions occur in the HVCs with the lowest H i column densities. This suggests that the total hydrogen column density distribution is smoother than indicated by the 21 cm emission maps.

Using arguments similar to the Si iii case, we can use the observed column densities of high ions (e.g., C iv and O vi) to infer the column densities of hot, collisionally ionized gas. The COS data on high ions (N v, C iv, O vi, Si iv, Si iii) are new or improved over earlier measurements. Our re-measurements of O vi HVC absorption are compared in Section 3.4 to previous FUSE studies (Sembach et al. 2003; Fox et al. 2004; CSG07). In many cases, the COS data have higher S/N, which we use to define the velocity range of HVC absorption. We re-measured O vi column densities using these ranges (see captions to Figures 4–7). The small differences with previous values (Sembach et al. 2003; Fox et al. 2004) are generally within the stated error bars. We adopt metallicities Z_C and Z_O of 10% solar and assume hot-gas ionization fractions $f_{\rm C\,\mathsc{iv}} = 0.3$ and $f_{\rm O\,\mathsc{vi}} = 0.2$. The observed C iv and O vi column densities toward three sight lines (Mrk 817, Mrk 290, Mrk 876) yield consistent "hot gas" column densities log N(H ii) =18.7–18.8, with C iv/O vi abundance ratios consistent, at N(C iv)/N(O vi) = 0.4–0.5. The fourth sight line (PG 1259+593) has the lowest C iv and O vi column densities, resulting in a lower inferred log N(H ii) =18.2–18.3 and a slightly lower ratio, N(C iv)/N(O vi) = 0.3.

The ionization ratios, C iv/O vi, N v/O vi, and Si iv/O vi, can be compared with observations of highly ionized HVCs (Fox et al. 2004, 2005, 2006; CSG05; CSG07; Indebetouw & Shull 2004b) and with models of various ionization processes (see Figure 1 of Indebetouw & Shull 2004a). In ratio plots of C iv/O vi and N v/O vi, the regions occupied by HVCs along our four sight lines are consistent with hot gas in conductive or radiatively cooling interfaces. The observed ratios are not consistent with collisional ionization equilibrium (CIE) or turbulent mixing layers (TMLs). For relative solar abundances (O:C:N = 1.00:0.55:0.14), the CIE models produce much higher N v/C iv ratios than observed; the obvious explanation is the lower nitrogen abundance observed in neutral HVC gas. Models of TMLs exhibit a wide range of predictions, owing to the assumptions and parameterizations that go into calculating the ionization and cooling. Early TML models (Slavin et al. 1993; Indebetouw & Shull 2004a) assumed mixing to an intermediate temperature, $\overline{T} \approx 10^5$ K, and relaxed to ionization equilibrium (Slavin et al. 1993; Indebetouw & Shull 2004a). Those models produce much higher C iv/O vi ratios than observed. More recent TML models (Kwak & Shelton 2010) that incorporate non-equilibrium ionization find regions with warm, radiatively cooled C iv, mixed with hotter gas, out of ionization equilibrium. These models find two to four times higher column densities in N v, C iv, and O vi abundances than predicted in CIE. However, their typical TML ratios are C iv/O vi = 1.5 (range 0.8–2.4) and N v/O vi (range 0.14–0.32), both higher than the COS observations, which have more O vi than predicted by the models. Models of non-equilibrium ionization (Gnat & Sternberg 2007) with time-dependent cooling (Z ≈ 0.1 Z_☉) find fair agreement with the COS-observed ionization ratios.

We observe a range of ratios, Si iv/C iv ≈ 0.15–0.36, along the four sight lines, typical of the previous studies (Fox et al. 2004, 2005; CSG05; CSG07). For relative solar abundances, (Si/C)_☉ = 0.11, photoionization models of HVCs find abundance ratios, N(Si iv)/N(C iv) >1, for values of photoionization parameter, log U = −3.0 ± 0.2, needed to explain the low ions and fit the ratios Si iv/Si iii/Si ii (Shull et al. 2009). These same models also underpredict the total column densities of high ions (C iv, Si iv, N v, O vi). Similarly, photoionization models with the observed range, N(Si iv)/N(C iv) = 0.1–0.3, predict N(Si iv)/N(Si ii) ≈1, much higher than observed. Therefore, it is likely that some C iv and Si iv comes from hot gas. The kinematic association of low and high ions in these HVCs requires a mixture of denser cloud cores of H i with extended warm photoionized gas and sheaths of much hotter gas, perhaps produced by bow shocks and turbulent or conductive interfaces between the HVC core and hot halo gas. However, detailed comparisons of observations and models are often complicated by the assumptions of relative solar abundances.

One of the reasons for renewed interest in Galactic HVCs is the possibility that the Local Group might contain considerable mass in virialized halos and a hot circumgalactic medium. Spitzer (1956) first suggested the existence of low-density coronal gas with T = 10⁶ K and n_e = 5 × 10⁻⁴ cm⁻³, extending 8 kpc above the Galactic plane and providing pressure confinement of observed high-latitude clouds. Kahn & Woltjer (1959) noted inconsistencies in Galactic stellar masses and Local Group dynamics (before the inference of dark matter) and suggested the existence of a substantial reservoir of low-density halo gas, with T = 5 × 10⁶ K and n_e = 1 × 10⁻⁴ cm⁻³. Direct probes of hot, low-density gas are difficult, owing to the n² dependence of its X-ray emission and the contamination of most signals by foreground electrons in the kpc-scale "Reynolds layer" (Reynolds 1991). Indirect probes of the halo density yield limits (n_e < 10⁻⁴ cm⁻³) from the effects of drag on orbits of the Magellanic Stream (Moore & Davis 1994). Similar limits (n_e < 3 × 10⁻⁵ cm⁻³) follow from ram-pressure stripping of Local Group dwarf galaxies (Blitz & Robishaw 2000). More recently, Heitsch & Putman (2009) used numerical simulations to suggest that infalling HVCs with H i masses less than 10^4.5 M_☉ may become fully ionized by Kelvin–Helmholtz instabilities within 10⁸ yr (∼10 kpc for typical HVC velocities of 100 km s⁻¹). All of these estimates depend critically on the assumed halo gas density and on dynamical interactions at the boundaries between the HVCs and the hot, low-density medium that confines them.

As a toy model for HVC cloud confinement, we consider spherical clouds with radius R, mass M, constant gas density ρ, and mean atomic mass μ = ρ/n_H = 1.23m_H (helium 25% by mass) in virial equilibrium confined by external pressure, P₀. For clumps of angular radius θ_deg (in degrees), we adopt R = (175pc)θ_deg at 10 kpc distance, n_H = (0.1 cm⁻³)n_0.1, and temperature T = (100 K)T₁₀₀. The virial theorem with confinement (Spitzer 1978) requires that

$$\begin{equation} 4 \pi R^3 \, P_0 = \frac{3MkT}{\mu } - \frac{3{\it GM}^2}{5R} \;. \end{equation} \tag{ 2 }$$

The confining pressure could arise from hot Galactic halo gas, which we scale to nominal values n_halo = (10⁻⁵ cm⁻³) n₋₅ and T_halo = (10⁶ K) T₆. Alternately, the HVCs could be confined by ram pressure, P_ram = ρ_haloV²_HVC, as they fall through the halo. For the assumed halo parameters, the total thermal pressure of fully ionized gas with n_He/n_H = 0.0823 is P/k ≈ 2.25n_HT = (22.5 cm⁻³ K)n₋₅ T₆. This pressure is consistent with inferences from various highly ionized HVCs (Sembach et al. 1999; CSG05), although such estimates are uncertain owing to assumptions in the ionization modeling. The halo density probably lies in the range n_e = (1–10) × 10⁻⁵ cm⁻³, with considerable variation over vertical distances 5–50 kpc above the Galactic plane.

In our model, the HVC clump masses in Complex C are M = (6.8 × 10⁴ M_☉)n_0.1θ³_deg, and the terms in the virial equation are all of comparable size,

$$\begin{eqnarray} (3MkT / \,\mu) &=& (2.7 \times 10^{48}\, {\rm erg}) \, n_{0.1}\, T_{100} \, \theta _{\rm deg}^3 \end{eqnarray} \tag{ 3 }$$

$$\begin{eqnarray} (3GM^2 / 5R) &=& (1.4 \times 10^{48}\, {\rm erg})\, n_{0.1}^2 \, \theta _{\rm deg}^5 \end{eqnarray} \tag{ 4 }$$

$$\begin{eqnarray} 4 \pi R^3 \, P_0 &=& (4.0 \times 10^{48}\, {\rm erg}) \, [n_{\rm halo}/10^{-5} {\rm cm}^{-3}] \, \theta _{\rm deg}^3 \;. \end{eqnarray} \tag{ 5 }$$

The thermal pressure of the halo, P_halo ≈ 1.6 × 10⁻¹⁵ erg cm⁻³, is comparable to the ram pressure on the HVCs. This not surprising, as these infalling clouds have low Mach numbers with respect to the hot gas. The HVC clumps may not be self-gravitating (no H₂ or stars have been detected). However, their observed properties place them near virial equilibrium (with T ⩽ 10³ K), and there may be no need to invoke dark matter for their confinement.

As the HVCs encounter the higher densities in the stratified lower halo, their outer portions may be torn apart and dissipated by interface instabilities. This would feed the Galactic halo rather than replenishing the reservoir of star formation in the disk. Portions of Complex C do appear to be clumping up, although the individual clump masses (θ_deg ⩾ 1) are likely above the 10^4.5 M_☉ threshold for survival (Heitsch & Putman 2009).

In this section, we compare the COS equivalent widths of several lines, previously measured by GHRS and STIS spectrographs (CSG03; CSG07; Richter et al. 2001). The derived column densities of various ions are compared in Tables 7–10. For some of the high ions (C iv, Si iv, Si iii, N v) our COS measurements are the only published data. Previous measurements of high-velocity Si iii toward Mrk 876 and PG 1259+593 were discussed in the HST surveys by Shull et al. (2009) and Collins et al. (2009).

• 1.  
  Mrk 817. The three unsaturated S ii lines at 1250, 1253, and 1259 Å should have equivalent widths in the ratio 1:2:3 based on their relative oscillator strengths. The weakest line (λ1250.6) was measured at W_λ = 27 ± 8 mÅ (COS) versus 16 ± 3 mÅ (GHRS). The stronger line (λ1253.8) was 19 ± 5 mÅ (COS) versus 28 ± 3 mÅ (GHRS). The strongest line (λ1259.5) with 42 ± 7 mÅ (COS) was not reported with GHRS. We base our S ii column density, log N = 14.29 ± 0.08, on the λ1259 measurement. For other lines we used Si iv λ1394, N v λ1238, Si ii λ1304 (consistent with λ1526), and Fe ii λ1608 (consistent with λ1145). We find good agreement between both lines in the doublets of C iv (λλ1548, 1551), N v (λλ1238, 1242), and Si iv (λλ1394, 1402). Fox et al. (2004) suggested that there may be no detectable high-velocity N v toward Mrk 817 because of intergalactic Lyα absorbers at z = 0.0189 and z = 0.0194. However, we believe we have detected N v with log N = 13.16^+0.12 _{− 0.15} at the −109 km s⁻¹ velocity of Complex C. As seen in Figure 8, the Lyα absorber at z = 0.0184 would appear at −42 km s⁻¹ in the N v rest frame, fairly well separated from the HVC. The ratios of N v to other high ions (C iv, Si iv, O vi) are typical of other sight lines. We use the weaker N i λ1199.5 line, since the other lines (1200.2 Å and 1200.7 Å) are blended with V_LSR = 0 absorption. Both Al ii λ1670.78 and Si iii λ1206.50 are mildly saturated, at equivalent widths of ∼400 mÅ and ∼350 mÅ, respectively. Our inferred column densities from AOD integration of N_a(v) accurately reflect this saturation, for line broadening of 20–30 km s⁻¹, somewhat higher than the Doppler parameter b = 11 km s⁻¹ inferred from O i and other low ions (CSG03). Re-measuring the O vi column density, we find log N = 14.05^+0.16 _{− 0.09} compared to 13.97^+0.10 _{− 0.13} from Fox et al. (2004) and 13.88 ± 0.20 (Sembach et al. 2003). The small differences in O vi column densities arise from the velocity range adopted for the HVC absorption. From the extent of UV absorption (S ii, Si ii, Fe ii, Al ii) in Figure 8, we integrate between V_LSR = −190 and −70 km s⁻¹, whereas Fox et al. (2004) used the interval from −160 to −80 km s⁻¹.
• 2.  
  Mrk 290. The S ii data from COS are superior to that from GHRS or STIS (CSG03; CSG07). The HVC components of λ1250.57 and λ1253.80 lie in the blue and red wings, respectively, of the broad Lyα emission line of the AGN (see Figure 5). To arrive at the column density, log N = 14.43 ± 0.08, we use the λ1259.51 line. For Si ii, both lines (λ1304 and λ1526) are saturated, yielding a lower limit on column density. We adopt the CoG value, log N = 14.9 ± 0.15, from CSG07. The column density from the mildly saturated Fe ii λ1608 is consistent with λ1145. We see agreement between both lines in the doublets of C iv (λλ1548, 1551), N v (λλ1238, 1242), and Si iv (λλ1394, 1402). We use the weaker line N i λ1199.5 line, since the other lines (1200.2 Å and 1200.7 Å) are blended with V_LSR = 0 absorption. Both Al ii λ1670.78 and Si iii λ1206.50 are mildly saturated, at equivalent widths of ∼340 mÅ and ∼345 mÅ, respectively. Our inferred column densities from AOD integration of N_a(v) accurately reflect this saturation, for line broadening of 20–30 km s⁻¹, somewhat higher than the Doppler parameter b = 16 km s⁻¹ inferred from O i and low ions (CSG07). We re-measured the O vi column density, finding reasonable agreement, log N = 14.10^+0.15 _{− 0.14} (between −175 and −75 km s⁻¹) compared to 14.23 ± 0.04 (CSG07 integrated between −165 and −75 km s⁻¹) and 14.20 ± 0.16 (Sembach et al. 2003).
• 3.  
  Mrk 876. The S ii data from COS are superior to that from STIS (CSG07), which only gave an upper limit, log N < 14.34. With COS, the HVC components of λ1259.51 and λ1253.80 yield essentially the same column density, log N = 14.25^+0.13 _{− 0.14}. For Si ii, both lines (λ1304 and λ1526) are mildly saturated, but yield a consistent column density, log N = 14.41^+0.16 _{− 0.09}. The column density from Fe ii λ1608, log N = 14.26^+0.17 _{− 0.10}, is consistent with that from λ1145. We find agreement between both lines in the doublets of C iv (λλ1548, 1551), N v (λλ1238, 1242), and Si iv (λλ1394, 1402). We use the weaker N i λ1199.5 line, since the other lines (1200.2 Å and 1200.7 Å) are blended with V_LSR = 0 absorption. Both Al ii λ1670.78 and Si iii λ1206.50 are mildly saturated, at equivalent widths of ∼280 mÅ and ∼270 mÅ, respectively. Our inferred column densities from AOD integration of N_a(v) accurately reflect this saturation, for line broadening of 20 km s⁻¹, somewhat higher than the Doppler parameter b = 16 km s⁻¹ inferred from O i and low ions (CSG03). Our previous HVC survey of Si iii (Shull et al. 2009) measured log N ⩾ 13.92 for this absorber, slightly higher than the value, 13.72^+0.24 _{− 0.14} measured here. We re-measured the O vi column density, finding moderate differences, but within stated errors, log N = 13.99^+0.36 _{− 0.14} compared to 14.20 ± 0.02 (CSG07), 14.05 ± 0.17 (Sembach et al. 2003), and 14.12^+0.11 _{− 0.13} (Fox et al. 2004). From the observed extent of the UV absorption (S ii, Si ii, Fe ii), we integrate O vi between V_LSR = −170 and −100 km s⁻¹, whereas Fox et al. (2004) used the interval from −220 to −100 km s⁻¹, and CSG07 used −210 to −75 km s⁻¹.
• 4.  
  PG 1259+593. The S ii data from COS are superior to that from GHRS (CSG03) and agree well with those from STIS (Sembach et al. 2004; Richter et al. 2001). Our column density, log N = 14.35^+0.15 _{− 0.10}, agrees with previous measurements. For Si ii, both lines (λ1304 and λ1526) are mildly saturated and yield a consistent column density. The column density from Fe ii λ1608 is consistent with λ1145. We find agreement between both lines in the doublets of C iv (λ1548, 1551), N v (λ1238, 1242), and Si iv (λ1394, 1402). Neither line in the N v doublet (λ1238, 1242) is seen to a limit log N < 12.99 (3σ). We use the weaker N i λ1199.5 line, since the other lines at 1200.2 Å and 1200.7 Å are blended with V_LSR = 0 absorption. Both Al ii λ1670.78 and Si iii λ1206.50 are mildly saturated, at equivalent widths of ∼215 m Å and ∼190 m Å, respectively. Our inferred column densities from AOD integration of N_a(v) accurately reflect this saturation, for line broadening of 20–30 km s⁻¹, somewhat higher than the Doppler parameter b = 10 km s⁻¹ inferred from O i and low ions (CSG03). Our previous HVC survey of Si iii (Shull et al. 2009) measured log N ⩾ 13.60 for this absorber, somewhat higher than the value, log N = 13.30^+0.24 _{− 0.16} measured here. We re-measured the O vi column density of this HVC, finding fair agreement, log N = 13.58^+0.18 _{− 0.12} compared to 13.71 ± 0.09 (Fox et al. 2004) and 13.72 ± 0.17 (Sembach et al. 2003). We integrate between V_LSR = −170 and −95 km s⁻¹, while Fox et al. (2004) used the interval from −160 to −80 km s⁻¹. For the O i column density in Complex C, we adopt the careful measurement, $\log N_{{\rm O\,\mathsc {i}}} = 15.85 \pm 0.15$ from Sembach et al. (2004), who fitted components as part of their study to find a deuterium ratio, D/H = (2.2 ± 0.7) × 10⁻⁵ in this HVC.