Iron group

In chemistry and physics, the iron group refers to elements that are in some way related to iron. These elements are relatively abundant both on Earth and elsewhere in the universe. The term is ambiguous in different contexts, and almost obsolete in chemistry.

Periodic table

The iron group in the periodic table referred to the elements iron, cobalt and nickel, that is the first row of group VIII (or VIIIB under the old numbering system.^[1]) In modern numbering, the iron group appears as three columns numbered group 8, 9 and 10. These metals, and the platinum group immediately below them, were set aside from the other elements as they show obvious similarities among themselves in their chemistry, but are not obviously related to any of the other groups.

The similarities in chemistry along what is now known as the first row of the transition metals were noted by Adolph Strecker in 1859.^[2] Newlands' "octaves" (1865) were harshly criticized for separating iron from cobalt and nickel.^[3] Mendeleev stressed that groups of "chemically analogous elements" could have similar atomic weights as well as atomic weights which increase by equal increments, both in his original 1869 paper^[4] and his 1889 Faraday Lecture.^[5]

Analytical chemistry

In the traditional methods of qualitative inorganic analysis, the iron group consists of those cations which

have soluble chlorides; and
are not precipitated as sulfides by hydrogen sulfide in acidic conditions;
are precipitated as hydroxides at around pH 10 (or less) in the presence of ammonia.

The main cations in the iron group are iron itself (Fe²⁺ and Fe³⁺), aluminium (Al³⁺) and chromium (Cr³⁺).^[6] If manganese is present in the sample, a small amount of hydrated manganese dioxide is often precipitated with the iron group hydroxides.^[6] Less common cations which are precipitated with the iron group include beryllium, titanium, zirconium, vanadium, uranium, thorium and cerium.^[7]

Astrophysics

The iron group in astrophysics is the group of elements from chromium to nickel which are substantially more abundant in the universe than those that come after them – or immediately before them – in order of atomic number.^[8] The study of the abundances of iron group elements relative to other elements in stars and supernovae allows the refinement of models of stellar evolution.

Abundances of the chemical elements in the Solar System. Note that the scale of the vertical axis is logarthimic. Hydrogen and helium are most common, from the Big Bang. The next three elements (Li, Be, B) are rare because they are poorly synthesized in the Big Bang and also in stars. The two general trends in the remaining stellar-produced elements are: (1) an alternation of abundance in elements as they have even or odd atomic numbers, and (2) a general decrease in abundance, as elements become heavier. The "iron peak" may be seen in the elements near iron as a secondary effect, increasing relative abundances of elements with nuclei most strongly bound.

The explanation for this relative abundance can be found in the process of nucleosynthesis in certain stars, specifically those of about 8–11 Solar masses. At the end of their lives, once other fuels have been exhausted, such stars can under a brief phase of "silicon burning".^[9] This involves the sequential addition of helium nuclei ⁴
₂He
(an "alpha process") to the heavier elements present in the star, starting from ²⁸
₁₄Si
:

²⁸ ₁₄Si	+	⁴ ₂He	→	³² ₁₆S
³² ₁₆S	+	⁴ ₂He	→	³⁶ ₁₈Ar
³⁶ ₁₈Ar	+	⁴ ₂He	→	⁴⁰ ₂₀Ca
⁴⁰ ₂₀Ca	+	⁴ ₂He	→	⁴⁴ ₂₂Ti ^{[note 1]}
⁴⁴ ₂₂Ti	+	⁴ ₂He	→	⁴⁸ ₂₄Cr
⁴⁸ ₂₄Cr	+	⁴ ₂He	→	⁵² ₂₆Fe
⁵² ₂₆Fe	+	⁴ ₂He	→	⁵⁶ ₂₈Ni

All of these nuclear reactions are exothermic, that is they release energy: the energy that is released partially offsets the gravitational contraction of the star. However, the series ends at ⁵⁶
₂₈Ni
, as the next reaction in the series,

⁵⁶
₂₈Ni

+

⁴
₂He

→

⁶⁰
₃₀Zn

is endothermic. With no further source of energy to support itself, the core of the star collapses on itself while the outer regions are blown off in a Type II supernova.^[9]

Nickel-56 is unstable with respect to beta decay, and the final stable product of silicon burning is ⁵⁶
₂₆Fe
.

⁵⁶ ₂₈Ni	→	⁵⁶ ₂₇Co	+	β⁺	t_1/2 = 6.075(10) d
⁵⁶ ₂₇Co	→	⁵⁶ ₂₆Fe	+	β⁺	t_1/2 = 77.233(27) d

Nuclide stability

	Nuclide mass^[10]	Mass defect^[11]	Binding energy per nucleon^[12]
⁶² ₂₈Ni	61.9283451(6) u	0.5700031(6) u	8.563872(10) MeV
⁵⁸ ₂₆Fe	57.9332756(8) u	0.5331899(8) u	8.563158(12) MeV
⁵⁶ ₂₆Fe	55.9349375(7) u	0.5141981(7) u	8.553080(12) MeV

It is often incorrectly assumed that iron-56 is the most stable of all the nuclides.^[8] This is not quite true: ⁶²
₂₈Ni
and ⁵⁸
₂₆Fe
have slightly higher binding energies per nucleon – that is, they are slightly more stable as nuclides – as can be seen from the table on the right.^[13] However, there are no rapid nucleosynthetic routes to these nuclides. There are several stable nuclides of elements from chromium to nickel around the top of the stability curve, accounting for their relative abundance in the universe. The nuclides which are not on the direct alpha-process pathway are formed by the so-called S-process, the capture of slow neutrons within the star.

Notes and references

Notes

^ In lighter stars, with less gravitational pressure, the alpha process is much slower and effectively stops at this stage as titanium-44 is unstable with respect to beta decay (t_1/2 = 60.0(11) years).